ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 In. Section 5 we explore the extent to which stars can be successfully assigned to the thin and thick discs by kinematic selection., In Section 3 we explore the extent to which stars can be successfully assigned to the thin and thick discs by kinematic selection.3The total mass of the MN. is amongst the most poorly known of all Galactic parameters.,The total mass of the MW is amongst the most poorly known of all Galactic parameters.4 Several estimates. give values around. 1.0.4.0107A. (see e.g... Wilkinson Evans 1999 and more references. therein).," Several estimates give values around $1.0-4.0\times 10^{12} \msun$ (see e.g., Wilkinson Evans 1999 and more references therein)."5 A more robust quantity is the mass within a certain large radius (50 kpc for example) constrained by the motions of satellite galaxies. elobular. clusters. the local escape velocity. of stars. etc.," A more robust quantity is the mass within a certain large radius (50 kpc for example) constrained by the motions of satellite galaxies, globular clusters, the local escape velocity of stars, etc."6 Using these constraints together with the observed rotation curve of the disc. Wochanek (1996) determined. the MW mass within 50 kpc.," Using these constraints together with the observed rotation curve of the disc, Kochanek (1996) determined the MW mass within 50 kpc."7 This mass implies a circular velocity Vu of 206n +(206 al 90% CLL).," This mass implies a circular velocity $_{50}$ of $206^{+10}_{-11}$ $206^{+22}_{-25}$, at $90\%$ C.L.)."8 A more recent determination of Vey by Wilkinson Lyvans (1999). constructed from the current data set of objects (27) with known distances and radial velocities at Calactocentric radii greater than 20 kpc. is consistent with the values ound by Ixochanek (1996).," A more recent determination of $V_{50}$ by Wilkinson Evans (1999), constructed from the current data set of objects (27) with known distances and radial velocities at Galactocentric radii greater than 20 kpc, is consistent with the values found by Kochanek (1996)."9 A so of 206n lis the observational constrain we use to fix the virial mass of the AIW model for à given ALALL, A $_{50}$ of $206^{+10}_{-11}$ is the observational constrain we use to fix the virial mass of the MW model for a given MAH.10" Since the MALI determines he halo density. profile. wwill be different for different. ALALIs: for the average NLALI. his mass is M,22.8«107M..."," Since the MAH determines the halo density profile, will be different for different MAHs; for the average MAH, this mass is $\mvir =2.8\times 10^{12}$."11 The spin parameter A mainly inlluences the stellar disc scale length and the surface density., The spin parameter $\lambda$ mainly influences the stellar disc scale length and the surface density.12 The stellar scale lenet= oof the MW. inferred. from observations in infrared. bands or estimated from dynamical constraints is typically 3.0+0.5 kpe (Sackett 1997). this becomes our second. Constraint.," The stellar scale length of the MW inferred from observations in infrared bands or estimated from dynamical constraints is typically $3.0\pm 0.5$ kpc (Sackett 1997), this becomes our second constraint."13 Llowever. see for example. Drimmel Spergel 2001. where values even smaller than these were found.," However, see for example Drimmel Spergel 2001 where values even smaller than these were found."14 For the average ALALL. the observed. central value of vvields A=0.02.," For the average MAH, the observed central value of yields $\lambda=0.02$."15" Once iis given. {y=Ma/M, is fixed by the MW. disc|bulge massAly."," Once is given, $\equiv \md/\mvir$ is fixed by the MW disc+bulge mass."16. Most barvons are in the disc|bulge svstem: 4.5.Lot? iin form of stars(AL... our third. constraint) and at leas ~O.6LO.107 iin form of gas (ALL) (see Table 1).," Most baryons are in the disc+bulge system: $4-5\times 10^{10}$ in form of stars, our third constraint) and at least $\sim 0.6-1.0 17\times 10^{10}$ in form of gas $_{\rm g}$ ) (see Table 1)."18" The estimated masses of the stellar halo (z4.6.10°AL... Haud Einasto 1989) anc of all MW satellite galaxies (~6910"" 3) represen only a minor contribution with respect to the disc|bulge svstem."," The estimated masses of the stellar halo $\approx194.6 \times 10^{9}$, Haud Einasto 1989) and of all MW satellite galaxies $\sim 6-9 \times 10^{9}$ ) represent only a minor contribution with respect to the disc+bulge system."20 Taking the average ΜΕΛΗ we obtain 50.091., Taking the average MAH we obtain =0.021.21 The inlluence of different \LALIs is felt upon the integra galaxy color index. gas fraction[... and scatter in the Tullv-Fisher relation.," The influence of different MAHs is felt upon the integral galaxy color index, gas fraction, and scatter in the Tully-Fisher relation."22 Unfortunately the NLALI cannot be describe by one-parameter and the color index and eas fraction are not accurately determined for the MW., Unfortunately the MAH cannot be described by one-parameter and the color index and gas fraction are not accurately determined for the MW.23 Thus. we are lef with a large range of possible ALALIs. all of which lead to models with the same Vso. aanel Ifor the MW.," Thus, we are left with a large range of possible MAHs, all of which lead to models with the same $V_{50}$, and for the MW."24 However. the SELL in the disc of the evolving galaxy. and particularly at the solar radius. is always highly sensitive to the input ALALL," However, the SFH in the disc of the evolving galaxy, and particularly at the solar radius, is always highly sensitive to the input MAH."25 The range of NLALIS allowed: by the structural constraints taken hence. defines a range of possible ος for the solar neighbourhood. with the most likely being that associated. to the average ALALI.," The range of MAHs allowed by the structural constraints taken hence defines a range of possible SFHs for the solar neighbourhood, with the most likely being that associated to the average MAH."26" A comparison with the inferred SELL of the solar neighbourhood. explored in δε, in fact shows the ALALL of our Galaxy to have been close to the average one for systems of the MN mass."," A comparison with the inferred SFH of the solar neighbourhood, explored in 4, in fact shows the MAH of our Galaxy to have been close to the average one for systems of the MW mass."27 The observational estimates of Vso. aancl geiven in Table 1 are thus the constraints we require in order for à model to be considered. representative. of he MW.," The observational estimates of $_{50}$, and given in Table 1 are thus the constraints we require in order for a model to be considered representative of the MW."28 For a given. ALALL the previous conditions. allow o determineMy... A andfy.," For a given MAH the previous conditions allow to determine, $\lambda$ and."29 When the average NAI js assumed we obtain —2.8LAL... A=0.02 and 200.021.," When the average MAH is assumed we obtain $=2.8 \times 10^{12}M_{\odot}$, $\lambda=0.02$ and 0.021."30 These conditions determine our MW central model. which represents a cosmological maximum: likelihood solution. given the structural parameters of our Galaxy ([icbucial moclel)," These conditions determine our MW central model, which represents a cosmological maximum likelihood solution, given the structural parameters of our Galaxy (fiducial model)."31 1t should be noted that the values of, It should be noted that the values of32A1795 (one of the better-filling clusters) and IHIvdra. A (the worst example).,A1795 (one of the better-fitting clusters) and Hydra A (the worst example).33 We see that the mixing moclel is generally consistent with the observations. (hough the \? is not always small.," We see that the mixing model is generally consistent with the observations, though the $\chi^2$ is not always small."34 Apart from providing a qualitative explanation for the observed density ancl temperature profiles of cluster gas. another major attraction of the thermal conduction model of ZNOS is (he fact that conduction helps to control the thermal instabilitv of the gas (INNO3).," Apart from providing a qualitative explanation for the observed density and temperature profiles of cluster gas, another major attraction of the thermal conduction model of ZN03 is the fact that conduction helps to control the thermal instability of the gas (KN03)."35 We have now repeated the global stability analvsis of INNO for the turbulent mixing model., We have now repeated the global stability analysis of KN03 for the turbulent mixing model.36" For A1795. with o,=0.011. we find that mixing suppresses the instability in all radial modes except the fundamental (nodeless) mode."," For A1795, with $\amix=0.011$, we find that mixing suppresses the instability in all radial modes except the fundamental (nodeless) mode."37 The growth time of the lone unstable mode is very much longer than the IIubble time., The growth time of the lone unstable mode is very much longer than the Hubble time.38" In the case of Ilvdra A. with oq,=0.021 (and thus a larger &ygi). we find that all modes. including the nodeless Iundamental mode. are stable."," In the case of Hydra A, with $\amix=0.021$ (and thus a larger $\kmix$ ), we find that all modes, including the nodeless fundamental mode, are stable."39 Thus. the equilibrium models described here are for all practical purposes stable.," Thus, the equilibrium models described here are for all practical purposes stable."40 Groups and clusters of galaxies form when primordial density perturbations in the universe grow. gravitationallv collapse. ancl merge together according to the stanclarcl hierarchical clustering scenario.," Groups and clusters of galaxies form when primordial density perturbations in the universe grow, gravitationally collapse, and merge together according to the standard hierarchical clustering scenario."41 The statistical properties of these collapsed systems contain many clues to the process of cosmic structure formation., The statistical properties of these collapsed systems contain many clues to the process of cosmic structure formation.42 Numerous X-ray studies have been published on the power-law scalings of the size. temperature. X-ray. Iuminosityv. mass. entropy. and gas mass fraction of galaxy clusters.," Numerous X-ray studies have been published on the power-law scalings of the size, temperature, X-ray luminosity, mass, entropy, and gas mass fraction of galaxy clusters."43 The best fit values of the power-law indices depend on the parüceular sample of clusters used. and on the specific methods emploved to estimate the mass and temperature., The best fit values of the power-law indices depend on the particular sample of clusters used and on the specific methods employed to estimate the mass and temperature.44 Nevertheless. (here is a broad consensus on the observed scalings as [unetions of temperature 7. as simmarized in the first three columns of Table 2..," Nevertheless, there is a broad consensus on the observed scalings as functions of temperature $T$, as summarized in the first three columns of Table \ref{tbl2}."45 A clear break in cluster properties is seen αἱ a characteristic temperature ~1—2 keV. sinall clusters ancl galaxy groups are found to have a relatively constant entropy., A clear break in cluster properties is seen at a characteristic temperature $\sim1-2$ keV. Small clusters and galaxy groups are found to have a relatively constant entropy.46" The prevailing explanations for this surprising ""entropy floor include (1) pre-heating of intracluster gas (Ixaiser1991:Evrard&Henry1991) via galactic winds (Loewenstein2000).. AGN (Valageas&Silk1999;WuXue2002:NateBRovehowdhlury2002).. and accretion shocks (Davidetal.2001:Tozzi&NormanDosSantosDoré2002).. (2) removal of cold. low-entropy gas via galaxy. formation in clusters (Bryan2000:Muanwongetal.2001:Wu&Xue2002:Daveetal. 2002).. and (3) both radiative cooling and supernova feedback (Voil&Dryan2001:Voitetal.2002)."," The prevailing explanations for this surprising “entropy floor” include (1) pre-heating of intracluster gas \citep{kai91,evr91} via galactic winds \citep{loe00}, AGN \citep{val99,wu02,nat02}, and accretion shocks \citep{dav01,toz01,dos02}, (2) removal of cold, low-entropy gas via galaxy formation in clusters \citep{bry00,mua01,wu02,dav02}, and (3) both radiative cooling and supernova feedback \citep{voit01,voit02}."47. None of these models includes the effects of thermal conduction or turbulent mixing., None of these models includes the effects of thermal conduction or turbulent mixing.48 When thermal conduction bv electrons is the dominant heating mechanism. the local," When thermal conduction by electrons is the dominant heating mechanism, the local"49nieasmiug the diverecuce of the displacement field) aud the deusitv field are prescuted in Section 3..,measuring the divergence of the displacement field and the density field are presented in Section \ref{sec:results}.50 Coucludiug remarks are given iu Section L.., Concluding remarks are given in Section \ref{sec:conclusion}.51 Iu this section. we compare the linear approximation of the contiuuitv equation to a logarithmic approximation.," In this section, we compare the linear approximation of the continuity equation to a logarithmic approximation."52 The large-scale dynamical evolution of structure is described by the Zeldovich approximation (Zoldovich 1970).. where x is the comoving Eulerian coordinate. q the Lagrangian coordinate. V the displacemeut. D the linear erowth function. aud * the exavitational potcutial.," The large-scale dynamical evolution of structure is described by the Zel'dovich approximation \citep{zel70}, where ${\bf x}$ is the comoving Eulerian coordinate, ${\bf q}$ the Lagrangian coordinate, $\bPsi$ the displacement, $D$ the linear growth function, and $\mathbf{\Phi}$ the gravitational potential."53 The linear Poisson equation relates the deusity coutrast. 6=(pp) /p. to the gravitational potential by 6=ΟνΦ. thus the divergence of the Lucar Zeldovich displacement can be related to the density contrast: The time derivative of the Zeldovich approximation iu expanding coordinates (r—ax. where r is the physical coordinate) allows us to relate this expression to the continuity equation.," The linear Poisson equation relates the density contrast, $\delta = (\rho - \bar{\rho})/\bar{\rho}$ , to the gravitational potential by $\delta = D\,\grad^2\mathbf{\Phi}$, thus the divergence of the linear Zel'dovich displacement can be related to the density contrast: The time derivative of the Zel'dovich approximation in expanding coordinates ${\bf r} = a{\bf x}$, where ${\bf r}$ is the physical coordinate) allows us to relate this expression to the continuity equation."54 We first note that the derivative consists of two compoucuts. r—6x|ax. where the first component describes the rate of expansion and the second the peculiar velocity of galaxies (or lass particles): v=exaW.," We first note that the derivative consists of two components, $\dot{\bf r} = \dot{a}{\bf x} + a\dot{\bf x}$, where the first component describes the rate of expansion and the second the peculiar velocity of galaxies (or mass particles): ${\bf v} = a\dot{\bf x} = a\dot\bPsi$."55 The continuity equation can i© Writtentt as ((sec.c.e..e..Peebles19801980.pLs To first in ὃ order this becomes thus we have d=Vo and have recovered Equation 2..," The continuity equation can be written as \citep[see, e.g.,][p 48]{pee80}56 To first in $\delta$ order this becomes thus we have $\dot\delta = -\grad\cdot\,\dot\bPsi$ and have recovered Equation \ref{eqn:divdel}. ."57 The contimutv equation (Eqn. 3)), The continuity equation (Eqn. \ref{eqn:continuity}) )58 may also be simplified by keeping the (1|4) together. resulting iu a logarithinic derivative of (1|à): This gives a logarithmic[m] expression for the divereeuce[m] of the Lagraugian[m] displacement. where the C term appears because ln(l|6) is not a zero-dnean-fBelkd. so we mst take ito account a non- iuteeration coustant.," may also be simplified by keeping the $(1+\delta)$ together, resulting in a logarithmic derivative of $(1+\delta)$: This gives a logarithmic expression for the divergence of the Lagrangian displacement, where the $C$ term appears because $\ln(1+\delta)$ is not a zero-mean-field, so we must take into account a non-zero integration constant."59 We fud that this constant is welL-approximated by the rather natural value of (utl| Aye., We find that this constant is well-approximated by the rather natural value of $\langle\ln(1+\delta)\rangle$ .60 We now may compare the linear relation. given bv Equ. 2..," We now may compare the linear relation, given by Eqn. \ref{eqn:divdel},"61 to the logarithmic relation. given by Equ. 6..," to the logarithmic relation, given by Eqn. \ref{eqn:divlnd}."62" We aeasure density field aud Wo.Ww im a 200 L tAIpe. 256° theparticle Gadget (Springeletal.2001) cold dark matter simulation with standard ACDAL cosinoloey (Qa,=0.3.04OF(Q7)"," We measure the density field and $\grad\cdot\bPsi$ in a 200 $^{-1}$ Mpc, $256^3$ particle Gadget \citep{spr01} cold dark matter simulation with standard $\Lambda$ CDM cosmology $\Omega_M = 0.3, \Omega_\Lambda = 0.7, h = 0.7$ )."63 We use three different methods: a erid-based Cloud-In-Cell (CIC) smoothing plus ai Fourier-space estimation of the divergence (Section 3.1)): an adaptive smoothing using a Smoothed Particle Wvdrodvuamics (SPIT) kernel. conibined with a Fourier-space divergence calculation (Section 3.2)): aud a geometrical Delaunay tesscllation estimation of both the density and the divergence (Section 3.3)).," We use three different methods: a grid-based Cloud-In-Cell (CIC) smoothing plus a Fourier-space estimation of the divergence (Section \ref{sec:cic}) ); an adaptive smoothing using a Smoothed Particle Hydrodynamics (SPH) kernel, combined with a Fourier-space divergence calculation (Section \ref{sec:sph}) ); and a geometrical Delaunay tessellation estimation of both the density and the divergence (Section \ref{sec:dtfe}) )."64 We find that cach of these methods achieves differing degrees of success based on the nature of the leneth-scales involved: the CIC inethod is an Eulerian. niass-weiehted scheme: the Delaunay tessellatiou is a Lagrangian. voluue-cweighted scheme: aud the adaptive mesh method is a hybrid. featuring a Lagraugian kerucl with interpolation outo an Eulerau exid.," We find that each of these methods achieves differing degrees of success based on the nature of the length-scales involved: the CIC method is an Eulerian, mass-weighted scheme; the Delaunay tessellation is a Lagrangian, volume-weighted scheme; and the adaptive mesh method is a hybrid, featuring a Lagrangian kernel with interpolation onto an Eulerian grid."65 We discuss cach in turn below., We discuss each in turn below.66 The CIC method smooths an arbitrary distribution of particles onto a reeularly-spaced exid bv placing a “cloud” haviug the volume of a grid cell around cach particle. so that a particle coutzibutes to the average of multiple nearby cells based on the fraction of its volume contained in these cells.," The CIC method smooths an arbitrary distribution of particles onto a regularly-spaced grid by placing a “cloud” having the volume of a grid cell around each particle, so that a particle contributes to the average of multiple nearby cells based on the fraction of its volume contained in these cells."67 Using this weighting scheme. we calculate the density aud displacement fields in real space for both 61? aud 128? cells. where the displaccuent W is the final (2=0) minus initial (2=19) particle positions.," Using this weighting scheme, we calculate the density and displacement fields in real space for both $64^3$ and $128^3$ cells, where the displacement $\bPsi$ is the final $z=0$ ) minus initial $z=49$ ) particle positions."68 Iu. the case of 128? cells at 2=0. there is oulv oue cell containiug zero particles: we sot its value to the average value iu the surrounding cells.," In the case of $128^3$ cells at $z=0$, there is only one cell containing zero particles; we set its value to the average value in the surrounding cells."69 To avoid cells with zero particles. this is the finest eril we use. eiven the resolution of the simulation.," To avoid cells with zero particles, this is the finest grid we use, given the resolution of the simulation."70" The divergence of the displacement can then casily be caleulated in Fourier space. where the derivatives become simple inultiples of hy,"," The divergence of the displacement can then easily be calculated in Fourier space, where the derivatives become simple multiples of $k$."71" We note here that BAO reconstruction mctlocds also calculate the linear Zel'dovich displacement iu Fourier space. which they iuultiplv by a smoothing function Sth): Win,=(ko,σκι for example. a Carssian snootling function would take the form οί)=(jD ?, "," We note here that BAO reconstruction methods also calculate the linear Zel'dovich displacement in Fourier space, which they multiply by a smoothing function $S(k)$: $\bPsi_{k,\rm{lin}} = i\,{\bf k}\,\delta_k\,S(k)/k^2$; for example, a Gaussian smoothing function would take the form $S(k) = e^{-k^2R^2/2}$ ."72We do not apply an extra smoothing fiction in Fourier space. but note that the size of the erid cells eives the effective. smoothing leneth of the both the deusity aud displacement fields.," We do not apply an extra smoothing function in Fourier space, but note that the size of the grid cells gives the effective smoothing length of the both the density and displacement fields."73 In our siunilations. a 61? cell exid las a leugth of 3.1 1Mpe aud a 128° cell grid has leugth 1.6L 1ALDpe. while the reconstruction mcthods have R= 5. 10. and 20 *Mpe. with 2—10 !Mpe performing the best (Eisensteinetal.," In our simulations, a $64^3$ cell grid has a length of 3.1 $^{-1}$ Mpc and a $128^3$ cell grid has length 1.6 $^{-1}$ Mpc, while the reconstruction methods have $R$ = 5, 10, and 20 $^{-1}$ Mpc, with $R$ =10 $^{-1}$ Mpc performing the best \citep{eisrec,noh09}."742007:Nolct 20049).. Iu Figure d we show a two-dineusional histogram of V(XP versus à aud lu(l|6) at four different redshifts. or a cell size of 1.6 |Mpe aud with a slope of -1 Xotted for reference.," In Figure \ref{fig:divcic4z} we show a two-dimensional histogram of $\grad\cdot\bPsi$ versus $\delta$ and $\ln(1+\delta)$ at four different redshifts, for a cell size of 1.6 $^{-1}$ Mpc and with a slope of -1 plotted for reference."75 The color scale is logarithiuic so hat the outlying cells can be seeu. though the furthest outliers (especially at low-:) extend bevoud the range of heplot.," The color scale is logarithmic so that the outlying cells can be seen, though the furthest outliers (especially at $z$ ) extend beyond the range of theplot."76 At := 7.the nonlinear clustering of matter as only begun for the initially highest-density peaks. so oth the linear and logarithmic relations have slopes near -] and are very tight.," At $z=7$ ,the nonlinear clustering of matter has only begun for the initially highest-density peaks, so both the linear and logarithmic relations have slopes near -1 and are very tight."77 ΑΕΙ=3 there is more scatter. and the linear relation beeius to deviate very slehtly roni a slope of -l for the majority of cells (shown iu," At $z=3$ there is more scatter, and the linear relation begins to deviate very slightly from a slope of -1 for the majority of cells (shown in"78sensitivity over that of CGRO (Lichtietal.1996:Ubertiui1996).,"sensitivity over that of CGRO \citep{lichti96,ubertini96}."79. In addition. with their much smaller ΕΟΝ. the INTEGRAL instruments. unlike COMPTEL. may be severely limited in terms of their total exposure to Cygnus X-1.," In addition, with their much smaller FoV, the INTEGRAL instruments, unlike COMPTEL, may be severely limited in terms of their total exposure to Cygnus X-1."80 COMPTEL may therefore provide the best available data on the MeV spectrum of Cyguus X-1 for many years to come., COMPTEL may therefore provide the best available data on the MeV spectrum of Cygnus X-1 for many years to come.81 The continued analysis of data from CGRO (only a fraction of the total COMPTEL data are used here) may help to further clarify the nature of this high euergy emission and perhaps enable us to more fully elucidate the physics oL the accretion process around stellar-1uass black holes., The continued analysis of data from CGRO (only a fraction of the total COMPTEL data are used here) may help to further clarify the nature of this high energy emission and perhaps enable us to more fully elucidate the physics of the accretion process around stellar-mass black holes.82 The authors would like to acknowledge Juri Poutanen aud. Andrzej Zdziarski for comments ou the original manuscript., The authors would like to acknowledge Juri Poutanen and Andrzej Zdziarski for comments on the original manuscript.83 Juri Poutanen kindly provided the NSPEC version of his νὰ thermal/non-thermal model (Poutanen&Svensson19906). that was used in this aulaysis., Juri Poutanen kindly provided the XSPEC version of his hybrid thermal/non-thermal model \citep{poutanen96} that was used in this anlaysis.84 The COMPTEL project is supported by NASA under contract NAS5S-26615. by the German government through DLR graut 50 C) 9096 8 aud by the Netherlands Oreanization for Scientilic Research NWO.," The COMPTEL project is supported by NASA under contract NAS5-26645, by the German government through DLR grant 50 Q 9096 8 and by the Netherlands Organization for Scientific Research NWO."85 This work has also been supported at UNH by the CGRO Guest Investigator Program uncer NASA evant NACS-7715., This work has also been supported at UNH by the CGRO Guest Investigator Program under NASA grant NAG5-7745.86"high energy of the cutoff in the Fermi//LAT spectrum of wwe favour a high-altitude location for gamma-ray production, leading to the higher value of the spin-down luminosity.","high energy of the cutoff in the /LAT spectrum of we favour a high-altitude location for gamma-ray production, leading to the higher value of the spin-down luminosity."87 Synchrotron emission will have no dependence on the orbital inclination because we are assuming isotropic particle distribution and a fully disordered magnetic field., Synchrotron emission will have no dependence on the orbital inclination because we are assuming isotropic particle distribution and a fully disordered magnetic field.88" For IC emission, on the other hand, there is a dependence with the angle 6 between the line of sight and the direction of propagation of the seed photons (i.e., the connecting line between the star and the compact object)."," For IC emission, on the other hand, there is a dependence with the angle $\theta$ between the line of sight and the direction of propagation of the seed photons (i.e., the connecting line between the star and the compact object)."89" Given the orbital configuration of 1),, 0 will not vary significantly with inclination for the phases around apastron, at which the source is detected in VHE."," Given the orbital configuration of , $\theta$ will not vary significantly with inclination for the phases around apastron, at which the source is detected in VHE."90" For yy absorption, on the other hand, the dependency is on the angle between the directions of the VHE photon and the seed photon integrating along the line of sight out of the system."," For $\gamma\gamma$ absorption, on the other hand, the dependency is on the angle between the directions of the VHE photon and the seed photon integrating along the line of sight out of the system."91 This means that absorption will show weaker orbital variability for low inclinations and stronger for high inclinations., This means that absorption will show weaker orbital variability for low inclinations and stronger for high inclinations.92" For the flux peak observed at $=0.62, the optical depth of yy absorption is practically independent of inclination at τγγ~0.2 for 500 GeV photons."," For the flux peak observed at $\phi=0.62$, the optical depth of $\gamma\gamma$ absorption is practically independent of inclination at $\tau_{\gamma\gamma}\approx0.2$ for 500 GeV photons."93" However, the second, wider outburst at phases 0.8—1.1, does have an absorption dependency with inclination, with t,,«1.4 for i=60° but τν~0.3 for i=15°."," However, the second, wider outburst at phases 0.8–1.1, does have an absorption dependency with inclination, with $\tau_{\gamma\gamma}\approx1.4$ for $i=60^\circ$ but $\tau_{\gamma\gamma}\approx0.3$ for $i=15^\circ$."94" The diminished absorption at low inclinations leads to a VHE flux higher than the observed values, particularly taking into account the upper limit measurement at phase 0.8."," The diminished absorption at low inclinations leads to a VHE flux higher than the observed values, particularly taking into account the upper limit measurement at phase 0.8."95" However, given the high uncertainties of the VHE nightly fluxes, the difference between the light curves computed for i=60° and i=15° is of less than lo."," However, given the high uncertainties of the VHE nightly fluxes, the difference between the light curves computed for $i=60^\circ$ and $i=15^\circ$ is of less than $1\sigma$."96" Even though we do not consider it significant given the present sensibilities, deeper VHE observations with firm detections at these phases could help to constrain the inclination of the orbit."," Even though we do not consider it significant given the present sensibilities, deeper VHE observations with firm detections at these phases could help to constrain the inclination of the orbit."97" The X-ray/VHE fluxes can be explained using a constant magnetic field of B—0.22 G. A constraint on the accelerator maximum energy can be put by requiring that the Larmor radius of the accelerated electron (rj,= E,/qBc) is contained within the accelerator size.", The X-ray/VHE fluxes can be explained using a constant magnetic field of $B=0.22$ G. A constraint on the accelerator maximum energy can be put by requiring that the Larmor radius of the accelerated electron $r_\mathrm{L}=E_\mathrm{e}/qBc$ ) is contained within the accelerator size.98" This translates to a maximum energy of Emax=300BGRi2 TTeV, where Bg is the magnetic field in Gauss and Rj;5 the accelerator radius in units of 10'? ccm."," This translates to a maximum energy of $E_\mathrm{max}=300\,B_\mathrm{G}R_{12}$ TeV, where $B_\mathrm{G}$ is the magnetic field in Gauss and $R_{12}$ the accelerator radius in units of $10^{12}$ cm."99" For the magnetic field we found and considering an accelerator size of the order of of the orbital separation, the maximum energy is well above the required 10 TeV and radiative losses dominate at high energies (see Sec. 3.3))."," For the magnetic field we found and considering an accelerator size of the order of of the orbital separation, the maximum energy is well above the required 10 TeV and radiative losses dominate at high energies (see Sec. \ref{sec:accel}) )."100" In the binary pulsar modelling approach of?,, the magnetic flux was considered variable along the orbit as a consequence of the variation in the distance of the wind-shock region to the pulsar."," In the binary pulsar modelling approach of, the magnetic flux was considered variable along the orbit as a consequence of the variation in the distance of the wind-shock region to the pulsar."101" In the shocked wind region, the magnetic field values obtained were as high as 10 G. On the other hand, in the microquasar approach used by?,, the magnetic field at the base of the jet was taken as 1 G. Most of these approaches use higher magnetic fields than the one we found to best reproduce the observational features."," In the shocked wind region, the magnetic field values obtained were as high as 10 G. On the other hand, in the microquasar approach used by, the magnetic field at the base of the jet was taken as 1 G. Most of these approaches use higher magnetic fields than the one we found to best reproduce the observational features."102 An explanation for this is likely the need in these models to generate both the pedestal and X-ray flux that we considered independently., An explanation for this is likely the need in these models to generate both the pedestal and X-ray flux that we considered independently.103 We presented a radiation model of a single leptonic population that can successfully describe the data obtained through a simultaneous X-ray/VHE campaign on pperformed by MAGIC in 2007., We presented a radiation model of a single leptonic population that can successfully describe the data obtained through a simultaneous X-ray/VHE campaign on performed by MAGIC in 2007.104 The observed X-ray/VHE correlation indicates both that the emission at these bands may come from a single particle population and that energy losses may be dominated by adiabatic losses., The observed X-ray/VHE correlation indicates both that the emission at these bands may come from a single particle population and that energy losses may be dominated by adiabatic losses.105" By assuming a constant magnetic field and injection along the orbit, we inferred the adiabatic loss rate from the X-ray light curve and the time dependent electron maximum energy as the balance between acceleration and energy loss time scales."," By assuming a constant magnetic field and injection along the orbit, we inferred the adiabatic loss rate from the X-ray light curve and the time dependent electron maximum energy as the balance between acceleration and energy loss time scales."106 We found that a quite efficient accelerator (7~ 10) is required to obtain the observed VHE spectra because of the fast adiabatic cooling., We found that a quite efficient accelerator $\eta\sim10$ ) is required to obtain the observed VHE spectra because of the fast adiabatic cooling.107" The injected electron spectrum is constant along the orbit and initially taken as a power law with a high-energy cutoff, but the Fermi//LAT HE spectrum poses a constraint on the hardness of the spectrum at lower energies, which we estimate has to be harder than o,~1.8 for electron energies below 4x10!! eV. At higher energies, on the other hand, an index of o,=2.1 matches the observed X-ray and VHE photon indices."," The injected electron spectrum is constant along the orbit and initially taken as a power law with a high-energy cutoff, but the /LAT HE spectrum poses a constraint on the hardness of the spectrum at lower energies, which we estimate has to be harder than $\alpha_\mathrm{e}\simeq1.8$ for electron energies below $4\times10^{11}$ eV. At higher energies, on the other hand, an index of $\alpha_\mathrm{e}=2.1$ matches the observed X-ray and VHE photon indices."108 The observed light curves are best reproduced using a magnetic field of B=0.22 G. The general picture shows that the observed emission in X-ray and VHE is fully compatible with originating in the same parent particle population under dominant adiabatic losses., The observed light curves are best reproduced using a magnetic field of $B=0.22$ G. The general picture shows that the observed emission in X-ray and VHE is fully compatible with originating in the same parent particle population under dominant adiabatic losses.109 The required luminosity budget in injected electrons for the computed X-ray and VHE emission is 2x10* erg/s. Both the miccroqquassar and binary pulsar scenarios are able to provide these levels of luminosity in accelerated electrons., The required luminosity budget in injected electrons for the computed X-ray and VHE emission is $2\times10^{35}$ erg/s. Both the sar and binary pulsar scenarios are able to provide these levels of luminosity in accelerated electrons.110" The GeV luminosity of ccan also give us hints on the available power: independently of whether the detected emission is magnetospheric or IC emission, the injected power should be z10°” erg/s. If it is magnetospheric emission, the high-altitude favoured scenario of gamma-ray production implies that the spin-down luminosity of the pulsar is E=3x10°” erg/s. This high spin-down luminosity poses problems for thebinary pulsar scenario because the pulsar wind would not be contained by the stellar wind, which would give rise to a completely different scenario from the one usually considered."," The GeV luminosity of can also give us hints on the available power: independently of whether the detected emission is magnetospheric or IC emission, the injected power should be $\ga 10^{37}$ erg/s. If it is magnetospheric emission, the high-altitude favoured scenario of gamma-ray production implies that the spin-down luminosity of the pulsar is $\dot{E}=3\times10^{37}$ erg/s. This high spin-down luminosity poses problems for thebinary pulsar scenario because the pulsar wind would not be contained by the stellar wind, which would give rise to a completely different scenario from the one usually considered."111take out rapid atinospheric fluctuatious.,take out rapid atmospheric fluctuations.112 The light curves preseuted here are self-calibrated tu pliase, The light curves presented here are self-calibrated in phase113"observations reported by n-Pintado et (1993), because they were collected in the most extended (A) configuration of the VLA.","observations reported by n-Pintado et (1993), because they were collected in the most extended (A) configuration of the VLA."114" Could our observations be dominated by extended. thermally excited gas. and those of n-Pintado et (1993) by denser, more compact masing regions?"," Could our observations be dominated by extended, thermally excited gas, and those of n-Pintado et (1993) by denser, more compact masing regions?"115 This appears not to be the case for the following reasons., This appears not to be the case for the following reasons.116" At a wavelength of 1.3 cm, the (E)VLA in the A configuration 1s sensitive to any structure more compact than about 2.5onds?."," At a wavelength of 1.3 cm, the (E)VLA in the A configuration is sensitive to any structure more compact than about 2.5."117. Our observations. on the other hand. show that the line emission detected here originates in a region more compact than about 1.5 arcseconds.," Our observations, on the other hand, show that the line emission detected here originates in a region more compact than about 1.5 arcseconds."118" Thus, there are no spatial scales contributing to the emission reported here and to which the observations by n-Pintado et (1993) were not sensitive."," Thus, there are no spatial scales contributing to the emission reported here and to which the observations by n-Pintado et (1993) were not sensitive."119" Moreover, at 1.3 cm, the size of MWC 349A is about x TTafoya et 2004: Rodrígguez et 22007: n-Pintado et 11993)."," Moreover, at 1.3 cm, the size of MWC 349A is about $\times$ Tafoya et 2004; guez et 2007; n-Pintado et 1993)."120 Since the angular resolution of the (E)VLA at 1.3 em in the A configuration is about071. MWC 349A only contains 4 independent resolution elements at that wavelength.," Since the angular resolution of the (E)VLA at 1.3 cm in the A configuration is about, MWC 349A only contains 4 independent resolution elements at that wavelength."121" Thus, the four spectra reported by n-Pintado et (1993) in their Figure 2 trace most of the extent of the source, and the continuum emission integrated over these four positions will be about 300 mJy (the total continuum emission estimated by n-Pintado et 11993)."," Thus, the four spectra reported by n-Pintado et (1993) in their Figure 2 trace most of the extent of the source, and the continuum emission integrated over these four positions will be about 300 mJy (the total continuum emission estimated by n-Pintado et 1993)."122" According to the spectra shown in Figure 2 of n-Pintado et ((1993), the mean intensity of the H66« line over the velocity range from —60 to --50 km | is about of the underlying continuum."," According to the spectra shown in Figure 2 of n-Pintado et (1993), the mean intensity of the $\alpha$ line over the velocity range from $-$ 60 to $+$ 50 km $^{-1}$ is about of the underlying continuum."123 The line intensity integrated over these 4 positions must therefore be about 60 mJy., The line intensity integrated over these 4 positions must therefore be about 60 mJy.124" As a consequence, the integrated line flux measured over any larger area should be at least 60 mJy in that velocity Yet. in our spectrum (corresponding to a x beam), the flux density is systematically smaller than 25 mJy: averaged over the —60 to 4-50 km ' velocity range. it is 18 mJy."," As a consequence, the integrated line flux measured over any larger area should be at least 60 mJy in that velocity Yet, in our spectrum (corresponding to a $\times$ beam), the flux density is systematically smaller than 25 mJy; averaged over the $-$ 60 to $+$ 50 km $^{-1}$ velocity range, it is 18 mJy."125 We conclude that n-Pintado et (1993) have overestimated the line intensity by a factor 3 to 4., We conclude that n-Pintado et (1993) have overestimated the line intensity by a factor 3 to 4.126" We can only speculate about the exact reasons which led n-Pintado et (1993) to over-estimate the H66o line intensity in MWC 349A. In principle, the lack of line-free channels in their observations should have led them to over-estimate the continuum flux (since the channels on the edge of the bandpass used to calculate the continuum still contained linc emission) and to under-estimate the line to continuum ratio (rather than the other way around)."," We can only speculate about the exact reasons which led n-Pintado et (1993) to over-estimate the $\alpha$ line intensity in MWC 349A. In principle, the lack of line-free channels in their observations should have led them to over-estimate the continuum flux (since the channels on the edge of the bandpass used to calculate the continuum still contained line emission) and to under-estimate the line to continuum ratio (rather than the other way around)."127sampling processes.,sampling processes.128 During the first 33.5 davs of science operations that followed the end of commissioning. the number of stars was increased {ο 156.000.," During the first 33.5 days of science operations that followed the end of commissioning, the number of stars was increased to 156,000."129 The volume of data for this nunber of stars observed everv 29.4 minutes is large. and the data from the 84 channels are diverse.," The volume of data for this number of stars observed every 29.4 minutes is large, and the data from the 84 channels are diverse."130 Thus it ds difficult to characterize (he entirety of the astrometric solution with a single number., Thus it is difficult to characterize the entirety of the astrometric solution with a single number.131 Again. much development in the astrometric processing is needed because every.Aepler star is important.," Again, much development in the astrometric processing is needed because every star is important."132 Although only preliminary measures of astrometric precision have been obtained. il is reassuring {ο see that the observed errors [ollow the prediction based on the flux of the stars involved.," Although only preliminary measures of astrometric precision have been obtained, it is reassuring to see that the observed errors follow the prediction based on the flux of the stars involved."133 Fig., Fig.134 1 shows the errors for stars in a single channel as a function of the measured.Ap. and demonstrates that the error rises as the SNR decreases.," 1 shows the errors for stars in a single channel as a function of the measured, and demonstrates that the error rises as the SNR decreases."135 A special case of the astrometric solutions described above can be computed in the vicinity of individual stars., A special case of the astrometric solutions described above can be computed in the vicinity of individual stars.136 Under the assumptions of small parallax aud adecuate removal of differential velocity aberration. (he equations for the apparent place of a star can be linearized.," Under the assumptions of small parallax and adequate removal of differential velocity aberration, the equations for the apparent place of a star can be linearized."137 Simple (vend analysis produces a robust estimator for Che mean position of a star. and residuals [rom each measurement are computed.," Simple trend analysis produces a robust estimator for the mean position of a star, and residuals from each measurement are computed."138 This enables astrometric processing to contribute to the understanding of theAepler stars., This enables astrometric processing to contribute to the understanding of the stars.139 As more filly discussed bv Batalhaetal.(2010).. what appears to be a single object can be two or more stars. and each can have photometric variability.," As more fully discussed by \citet{bah10}, what appears to be a single object can be two or more stars, and each can have photometric variability."140 Such astrations can mimic the photometric properties of a transiting planet. ancl adding astrometry to (he velGine procedure can assist in the confirmation or denial process.," Such astrations can mimic the photometric properties of a transiting planet, and adding astrometry to the vetting procedure can assist in the confirmation or denial process."141 Fie., Fig.142 2 shows astrometric residuals for (wo stars., 2 shows astrometric residuals for two stars.143 The upper star i$ a blend of a variable star and one or more constant stars while the lower shows residuals that are tvpical for a bright. constant star.," The upper star is a blend of a variable star and one or more constant stars while the lower shows residuals that are typical for a bright, constant star."144 The astrometric amplitude of the variable star is huge - almost 0.02 pixels., The astrometric amplitude of the variable star is huge - almost 0.02 pixels.145"O, agrees well for all (τους kinematic methods aud all subgroups.",$\Theta_0$ agrees well for all three kinematic methods and all subgroups.146 The inner 5 kpe region has a different rotation axis than the outer regions. demonstrated clearly in all three binning methods.," The inner 5 kpc region has a different rotation axis than the outer regions, demonstrated clearly in all three binning methods."147" The innermost bin vields weighted averages of OQ,=369424. O,=252:55"". and Q,=352cI8"". all of which are equal within their uncertainties."," The innermost bin yields weighted averages of $\Theta_o = 369\pm24 ^o$, $\Theta_o148= 25\pm55 ^o$, and $\Theta_o = 352\pm18 ^o$, all of which are equal within their uncertainties."149" Devond 5 kpc. the rotation axes lor all three populations are in even closer agreement. wilh averages of O,=189c6°. Q,—199ctτὸ and O,=196x1"" lor the entire population. the metal-poor subpopulation. and the metal-xrich subpopulation. respectively."," Beyond 5 kpc, the rotation axes for all three populations are in even closer agreement, with averages of $\Theta_o = 189\pm6 ^o$, $\Theta_o = 199\pm7 ^o$, and $\Theta_o = 196\pm7 ^o$ for the entire population, the metal-poor subpopulation, and the metal-rich subpopulation, respectively."150 The position angle of (he photometric major axis of NGC 5128 is O=35° and 215° east of north and the photometric munor axis is Ο=119° and 299° east of north (Dufouretal.1979)., The position angle of the photometric major axis of NGC 5128 is $\Theta = 35^o$ and $215^o$ east of north and the photometric minor axis is $\Theta = 119^o$ and $299^o$ east of north \citep{dufour79}.151. It appears the GCS is rotating about an axis similar to the photometric major axis for the fill extent of the galaxy. with a possible axial twist or counterrotation within 5 kpe.," It appears the GCS is rotating about an axis similar to the photometric major axis for the full extent of the galaxy, with a possible axial twist or counterrotation within 5 kpc."152 Fieures 9 10 show the velocity dispersion for (he entire population ancl [or the metal-poor and metal-rich subpopulations.," Figures \ref{fig:veldisp_final} \ref{fig:veldisp_metal}153 show the velocity dispersion for the entire population and for the metal-poor and metal-rich subpopulations."154" Our results for e, show no significant differencesbetween the metallicity subpopulations.", Our results for $\sigma_v$ show no significant differencesbetween the metallicity subpopulations.155" All three show a relatively flat. velocity dispersion (c,=O44. σι=LITAG. and σι=11146 kms ! within 15 kpe of the center of NGC! 5128 for the entire population anc for the metal-poor ancl metal-rich subpopulations. respectively)."," All three show a relatively flat velocity dispersion $\sigma_v = 119\pm4$, $\sigma_v = 117\pm6$, and $\sigma_v = 111\pm6$ km $^{-1}$ within 15 kpc of the center of NGC 5128 for the entire population and for the metal-poor and metal-rich subpopulations, respectively)."156 These results match the previous study of NGC 5128 by Peng. (2004c).. whose determined velocity dispersion for the GCs within 20 kpc ranged between τὸ and 150 km !..," These results match the previous study of NGC 5128 by \cite{pff04II}, whose determined velocity dispersion for the GCs within 20 kpc ranged between 75 and 150 km $^{-1}$."157" At a larger radius. we find that o, then slowly increases to a,>150 km | towards the outer regions of the halo for all populations."," At a larger radius, we find that $\sigma_v$ then slowly increases to $\sigma_v > 150$ km $^{-1}$ towards the outer regions of the halo for all populations."158 The velocity dispersion of the metalrich GCs. inlerestinely. appears higher than that of the metal-poor GC's in the outer regions (although still consistent within the determined uncertainties).," The velocity dispersion of the metal-rich GCs, interestingly, appears higher than that of the metal-poor GCs in the outer regions (although still consistent within the determined uncertainties)."159 [a most previous studies. the velocity dispersion of the metal-poor GCs usually appears higher (han that of (he metal-rich GC's. if there is a notable velocity dispersion difference between the subpopulations (seethestudiesofCotéetal.2003:Richtlerοἱ2004.asexamples).," In most previous studies, the velocity dispersion of the metal-poor GCs usually appears higher than that of the metal-rich GCs, if there is a notable velocity dispersion difference between the subpopulations \cite[see the studies of][as examples]{cote03,richtler04}."160 To explore the cause of the disüncet rise past 15 kpe a bit further. we have plotted the actual velocity histogramsex in Figuree 11. for the metal-poor and metal-rich subgroups.e subcdivided huther into inner (2<15 kpe) aud outer (22>15 kpc) regions.," To explore the cause of the distinct rise past 15 kpc a bit further, we have plotted the actual velocity histograms in Figure \ref{fig:vf_histo} for the metal-poor and metal-rich subgroups, subdivided further into inner $R < 15$ kpc) and outer $R > 15$ kpc) regions."161 In the inner 15 kpc. both samples show histograms strongly peaked near v;=0 and with at least roughly Gaussian falloff to both high and low velocities.," In the inner 15 kpc, both samples show histograms strongly peaked near $v_f = 0$ and with at least roughly Gaussian falloff to both high and low velocities."162 By contrast. the histograms for the outer regions (15—50 kpc) are noliceably flatler. so that the clusters with lareer velocity residuals have relatively more importance (ο the formal value of σι.," By contrast, the histograms for the outer regions $15-50$ kpc) are noticeably flatter, so that the clusters with larger velocity residuals have relatively more importance to the formal value of $\sigma_v$ ."163 Nominally. the [latter shape of the velocity distribution would mean that the outer-halo clusters display anisotropy. in," Nominally, the flatter shape of the velocity distribution would mean that the outer-halo clusters display anisotropy in"164(second acoustic peak) etc.,(second acoustic peak) etc.165 The squeezed configuration in turn is always negative as the smallest multipole has negative transfer function via SW., The squeezed configuration in turn is always negative as the smallest multipole has negative transfer function via SW.166 The folded configuration has relatively large negative values when the two smallest multipoles are in the first acoustic peak while the biggest multipole is in the second acoustic peak (P~900—1200>P/3~300—400)., The folded configuration has relatively large negative values when the two smallest multipoles are in the first acoustic peak while the biggest multipole is in the second acoustic peak $P\sim 900-1200 \Rightarrow P/3\sim300-400$ ).167 For larger perimeters the structure becomes complex as several acoustic peaks intervene., For larger perimeters the structure becomes complex as several acoustic peaks intervene.168" In this section we develop a prescription which permits us to predict the (bi)spectrum of point sources, starting from the simplest case of a single randomly-distributed population to the case of multiple clustered populations."," In this section we develop a prescription which permits us to predict the (bi)spectrum of point sources, starting from the simplest case of a single randomly-distributed population to the case of multiple clustered populations."169 This is a situation that is encountered in current and future CMB analyses., This is a situation that is encountered in current and future CMB analyses.170" Indeed until recently, CMB experiments have focussed on frequencies where unclustered radio sources are the only dominant kind of point sources, but the CMB is also non-negligible at higher frequencies where an independent population of dusty galaxies becomes important together with the SZ signal of clusters."," Indeed until recently, CMB experiments have focussed on frequencies where unclustered radio sources are the only dominant kind of point sources, but the CMB is also non-negligible at higher frequencies where an independent population of dusty galaxies becomes important together with the SZ signal of clusters."171 This is of particular relevance for Planck which has a large frequency range covering both populations., This is of particular relevance for Planck which has a large frequency range covering both populations.172" A source with flux S enclosed in a pixel with solid angle Qpix yields a temperature variation AT=kygaz where kp=BB) , B(v,T) is the black-body spectrum and Τον is the CMB mean temperature."," A source with flux $S$ enclosed in a pixel with solid angle $\Omega_{\mathrm{pix}}$ yields a temperature variation $\Delta T =173k_\nu \frac{S}{\Omega_{\mathrm{pix}}}$, where $k_\nu =174\left.\frac{\partial B(\nu,T)}{\partial T}\right|_{T_\mathrm{CMB}}$ , $B(\nu,T)$ is the black-body spectrum and $T_\mathrm{CMB}$ is the CMB mean temperature."175" As shown in Appendix ?? the power spectrum of a source population is given by: The discreteness of the sources produces a constant-term spectrum Cj""** which is usually named ‘Poissonian’ because the number of unclustered point sources is driven a priori by Poisson statistics (7)..", As shown in Appendix \ref{appendix:shot} the power spectrum of a source population is given by: The discreteness of the sources produces a constant-term spectrum $C_\ell^{\mathrm{shot}}$ which is usually named `Poissonian' because the number of unclustered point sources is driven a priori by Poisson statistics \citep{Sehgal2010}.176" The shot-noise term reads: where an is the differential number counts of sources and Scut is the detection limit, i.e sources with S>Sc are detected and masked, the rest being unresolved."," The shot-noise term reads: where $\frac{\mathrm{d}n}{\mathrm{d}S}$ is the differential number counts of sources and $S_\mathrm{cut}$ is the detection limit, i.e sources with $S>S_\mathrm{cut}$ are detected and masked, the rest being unresolved."177" The discreteness property of sources, when computing the three-point correlation function, yields a statistically isotropic angular bispectrum constant with 6: for £;zz0."," The discreteness property of sources, when computing the three-point correlation function, yields a statistically isotropic angular bispectrum constant with $\ell$: for $\ell_i \ne 0$."178" Equations (18)), (19)) and (20)) are derived in more detail in Appendix ??.."," Equations \ref{eq:PS_cl}) ), \ref{shotspec}) ) and \ref{shotbisp}) ) are derived in more detail in Appendix \ref{appendix:shot}."179 The case of sources randomly and independently distributed on the sky is that of the radio sources at CMB frequencies., The case of sources randomly and independently distributed on the sky is that of the radio sources at CMB frequencies.180 The correlation vanishes and the total (bi)spectrum isequal to the shot-noise (bi)spectrum., The correlation vanishes and the total (bi)spectrum isequal to the shot-noise (bi)spectrum.181 The distribution of the sources is that of a white-noise entirely, The distribution of the sources is that of a white-noise entirely182AL=ὃν98 is then derived using the finite cillerencing method.,$A_{\rm s}={\partial{\vec{\beta}_N}}/{\partial{\vec{\theta}_1}}$ is then derived using the finite differencing method.183 Images of a given source position are effectively found using the Newton-Raphson iteration method., Images of a given source position are effectively found using the Newton-Raphson iteration method.184 In order to calculate PCSRow|A@#£2.5°) the probability distribution in the f?A plane for individual lensing cones. we generate 10.000~20.000 cusp sources whose close triple images have image opening angles A@<1207.," In order to calculate $P(\geqslant \Rcusp | \Delta\theta185\pm2.5^{\circ})$ – the probability distribution in the $\Rcusp-\Delta\theta$ plane for individual lensing cones, we generate $10,000\sim20,000$ cusp sources whose close triple images have image opening angles $\Delta\theta \leqslant 120^{\circ}$ ."186 We have also caleulated DUquem) for cases with 4890 as an overall estimate for cusp-caustic violations to compare with our previous work. in which onlv cases with Adx90° were examined for violations (caused by intrinsic substructures within the main lens)," We have also calculated $P^{90}(\Rcusp^{0.187})$ for cases with $\Delta\theta187\leqslant 90^{\circ}$ as an overall estimate for cusp-caustic violations to compare with our previous work, in which only cases with $\Delta\theta \leqslant 90^{\circ}$ were examined for violations (caused by intrinsic substructures within the main lens)."188 To derive average violation probabilities over all sight ines.. we weight. P(BassN9E.2.5)B and PoLUm){NY of each lensing cone by the quacdruple-image lensing cross-section in the source plane (simply the fractional area within he tangential caustic).," To derive average violation probabilities over all sight lines, we weight $P(\geqslant \Rcusp | \Delta\theta \pm2.5^{\circ})$ and $P^{90}(\Rcusp^{0.187})$ of each lensing cone by the quadruple-image lensing cross-section in the source plane (simply the fractional area within the tangential caustic)."189 We do not account for magnification jas among the cusp sources., We do not account for magnification bias among the cusp sources.190 Fig., Fig.191 9. (upper panels) and Table 2 show that if »erturbing structures are haloes (and subhalocs) distributed outside the main lens along the line of sight as in the Millennium-HLE. simulation. where such haloes are resolved o 10h.AL... they cause a non-negligible amount of cusp violations. comparable to those due to the substructures in he lens itself," \ref{fig:losvsmainSISNFW} (upper panels) and Table \ref{tab:MSIIlosRcusp} show that if perturbing structures are haloes (and subhaloes) distributed outside the main lens along the line of sight as in the Millennium-II simulation, where such haloes are resolved to $10^8 h^{-1}M_{\odot}$, they cause a non-negligible amount of cusp violations, comparable to those due to the substructures in the lens itself."192 However. the violation pattern (as a function of A@) depends strongly on the density. profiles applied: to aloes projected near the centre of the line of sight.," However, the violation pattern (as a function of $\Delta\theta$ ) depends strongly on the density profiles applied to haloes projected near the centre of the line of sight."193 We have also investigated the effect from line-of-sight haloes more massive than LOTAL. which are most Likely to retain a significant fraction of barvons in their cark matter potential wells.," We have also investigated the effect from line-of-sight haloes more massive than $10^{10} h^{-1}M_{\odot}$, which are most likely to retain a significant fraction of baryons in their dark matter potential wells."194 The chance of finding at least one of these massive secondary Lenses intercepting a strong-lcnsing sight line (Le. projected. within a typical Einstein radius of 1” around the line centre. out to a redshift. of 2) is about1o%.," The chance of finding at least one of these massive secondary lenses intercepting a strong-lensing sight line (i.e. projected within a typical Einstein radius of $\arcsec$ around the line centre, out to a redshift of 2) is about."195. Depending on their density profiles. compact haloes (c.g. singular isothermal spheres) could generate severe astrometry anomalies with a probability of a few percent. while haloes with shallower inner profiles could not.," Depending on their density profiles, compact haloes (e.g. singular isothermal spheres) could generate severe astrometry anomalies with a probability of a few percent, while haloes with shallower inner profiles could not."196 Fig., Fig.197 10. presents the peculiar image configurations for an example sight line., \ref{fig:LC2CaseStudy} presents the peculiar image configurations for an example sight line.198 Four. six and eight images (excluding the central image) are produced when the source is located at different positions with respect to two sets of tangential caustics.," Four, six and eight images (excluding the central image) are produced when the source is located at different positions with respect to two sets of tangential caustics."199 In this particular case. the second caustic is produced by a perturbing halo of 2LOMA7.. modelled. as a truncated. singular isothermal sphere. projected near the centre ofthe main lens.," In this particular case, the second caustic is produced by a perturbing halo of $2\times10^{10} M_{\odot}$, modelled as a truncated singular isothermal sphere, projected near the centre of the main lens."200 Such peculiar image astrometry has already been proposed and used to constrain density profiles of massive intergalactic objects (e.g... 2?))," Such peculiar image astrometry has already been proposed and used to constrain density profiles of massive intergalactic objects (e.g., \citealt{Wyithe2001,Wilkinson2001}) )."201" ‘To investigate the ellect of substructures inside halocs along the line of sight. we have exeluded: all substructures from our Alillennitm-ll lensing cones ancl caleulated violations due to ""smooth"" line-of-sight haloes alone."," To investigate the effect of substructures inside haloes along the line of sight, we have excluded all substructures from our Millennium-II lensing cones and calculated violations due to “smooth” line-of-sight haloes alone."202 Table 2 lists violation probabilities in this case (for clilferent halo density profiles).," Table \ref{tab:MSIIlosRcusp}203 lists violation probabilities in this case (for different halo density profiles)."204 The relevance of subhaloes to lensing ILux-ratio anomalies strongly depends on their assumed density profiles., The relevance of subhaloes to lensing flux-ratio anomalies strongly depends on their assumed density profiles.205 Comparison of the first ancl second. rows of Table 2 shows that for m>107.TAL... NEW-like substructures within linc-of-sight haloes are responsible for causing a few percent of the cusp-caustic violations.," Comparison of the first and second rows of Table 2 shows that for $m > 10^{8}206h^{-1}M_{\odot}$, NFW-like substructures within line-of-sight haloes are responsible for causing a few percent of the cusp-caustic violations."207 We have separated line-of-5ight haloes that are distributed in front of ancl behind the main-lens plane (2;=0.6)., We have separated line-of-sight haloes that are distributed in front of and behind the main-lens plane $z_d=0.6$ ).208 The violation probabilities of these two groups are listed in Table 2.. calculated excluding their subhaloes.," The violation probabilities of these two groups are listed in Table \ref{tab:MSIIlosRcusp}, calculated excluding their subhaloes."209 A higher violation probability is foundcaused by haloes in the background than in the foreground. as more haloes intercept the light. ravs behind the main lens plane. given à typical lensing geometry (cu0.6 and ος=2).," A higher violation probability is foundcaused by haloes in the background than in the foreground, as more haloes intercept the light rays behind the main lens plane, given a typical lensing geometry $z_d=0.6$ and $z_s=2$ )."210 Ht is interesting to notice that violations from the foreground and the background roughly acid up to the total violations due to haloes along the entire line of sight (second row of Table 2))., It is interesting to notice that violations from the foreground and the background roughly add up to the total violations due to haloes along the entire line of sight (second row of Table \ref{tab:MSIIlosRcusp}) ).211 Phe ratio between violations from the foreground and from the entire sight line is close to 2:5. which is the ratio between the comoving racial distances for z;—0.6 and ον=2.," The ratio between violations from the foreground and from the entire sight line is close to 2:5, which is the ratio between the comoving radial distances for $z_d=0.6$ and $z_s=2$."212 The Millennium-L simulation has a limited mass resolution., The Millennium-II simulation has a limited mass resolution.213 To investigate the mass dependence of the violation pattern below the halo mass of 1075.ΑΙ. we have used a Monte- method. to generate intergalactic halo populations with masses 1095AJ.moo<107h181. (see 87 [or discussion. on adopting 105bTAL. as the lower mass limit).," To investigate the mass dependence of the violation pattern below the halo mass of $10^{8} h^{-1}M_{\odot}$, we have used a Monte-Carlo method to generate intergalactic halo populations with masses $10^{6}214h^{-1}M_{\odot} \leqslant m < 10^{12} h^{-1}M_{\odot}$ (see $\S$ 7 for discussion on adopting $10^{6} h^{-1}M_{\odot}$ as the lower mass limit)."215 These haloes are drawn from the Sheth-Tormen mass function (23) generated with the code provided by 2.., These haloes are drawn from the Sheth-Tormen mass function \citealt{STMassFunction2002}) ) generated with the code provided by \citet{ReedMassFunction2007}.216 We have randomly. generated 200 lensing cones out to 2.=2. each of which contains a main lensing halo mocelled as an isothermal ellipsoid at redshift 2;=0.6.," We have randomly generated 200 lensing cones out to $z_s=2$, each of which contains a main lensing halo modelled as an isothermal ellipsoid at redshift $z_d=0.6$."217" The lensing strength bse is fixed to be 0.84"". the same as the mean be of the main lenses in the selected sample of the Millennium-LL lensing cones."," The lensing strength $b_{\rm SIE}$ is fixed to be $0.84\arcsec$, the same as the mean $b_{\rm SIE}$ of the main lenses in the selected sample of the Millennium-II lensing cones."218 The axis ratio d;=0.8. core radius Sy1.05” and an orientation angle of 0.25; are also taken to be he same for all main lenses.," The axis ratio $q_3=0.8$, core radius $S_0=0.05\arcsec$ and an orientation angle of $\pi$ are also taken to be the same for all main lenses."219 1n each realization of the lensing cone. line-of-sight halo »ositions are randomly generated. with number densities as eiven by the Sheth-Tormen mass function at the redshifts of he GO lens planes used for the Millennium-LH lensing cones.," In each realization of the lensing cone, line-of-sight halo positions are randomly generated, with number densities as given by the Sheth-Tormen mass function at the redshifts of the 60 lens planes used for the Millennium-II lensing cones."220 Table 3. lists the mean surface number densitiesof projected iaoes in cüllerent mass decades. averaged over 200 lensing cones!.," Table \ref{tab:MChaloSurfaceNumber} lists the mean surface number densitiesof projected haloes in different mass decades, averaged over 200 lensing ."221". Llaloes projected. within the 50"". 50""-cone are saved. and those projected within the central 5”. 5""-region are modelled with truncated singular isothermal spheres ancl truncated NEW) profiles (using both MOS. ancl 101- concentration-mass relations)."," Haloes projected within the $50\arcsec\times22250\arcsec$ -cone are saved, and those projected within the central $5\arcsec\times 5\arcsec$ -region are modelled with truncated singular isothermal spheres and truncated NFW profiles (using both M08 and B01-M05 concentration-mass relations)."223 ''hose further out are, Those further out are224"An increase in metallicity could remove the cliscrepaney by reducing the mass-to-near-infraved-light ratios in Another important. factor. particularly for the HUC 5.Sjm band. is the possible contamination bv emission from hot dust. specifically. emission in the 3.3. 6.2 ancl 7.7, PAIL features €2?7)..","An increase in metallicity could remove the discrepancy by reducing the mass-to-near-infrared-light ratios in Another important factor, particularly for the IRAC $5.8\rm{\mu m}$ band, is the possible contamination by emission from hot dust, specifically, emission in the 3.3, 6.2 and $\rm{\mu m}$ PAH features \citep{Draine2007a, Draine2007b, DaCunha2008}."225 Such emission is not included in the models but may. well be significant in the real Iow-redshift galaxies., Such emission is not included in the models but may well be significant in the real low-redshift galaxies.226 In Fig., In Fig.227" 3. we compare predictions from our mocdels to the observed. photometric redshift. distributions of galaxy samples selected above A, and IRAC 5.5jun. apparen magnitude limits.", \ref{fig:zdist} we compare predictions from our models to the observed photometric redshift distributions of galaxy samples selected above $K_s$ and IRAC $5.8\rm{\mu m}$ apparent magnitude limits.228 As for the number counts. the solid. rec and dashed: blue lines represent predictions based on the ? and ?/ stellar population synthesis models.," As for the number counts, the solid red and dashed blue lines represent predictions based on the \citet{Maraston2005} and \citet{Bruzual2003} stellar population synthesis models."229 Filled: regions show the expected. le scatter among fields with an area of 100 arcmin?., Filled regions show the expected $\sigma$ scatter among fields with an area of 100 $^2$.230" In the left panel the number of galaxies per unit area and redshift is plotted for samples with A,«21.8 and AG« 23.3. while the right panel gives similar results bu or samples with LRAC ο.δι apparent magnitude brighter han 21.8 (and A,« 23.0)."," In the left panel the number of galaxies per unit area and redshift is plotted for samples with $K_s<21.8$ and $K_s<23.3$ , while the right panel gives similar results but for samples with IRAC $5.8\rm{\mu m}$ apparent magnitude brighter than 21.8 (and $K_s<23.0$ )."231 Our theoretical predictions are compared with data from two public catalogues: the FIREWORKS data for the GOODS-CDE (?) and. the -EWEIRAL Medium. Band Survey data for the COSMOS and ALGIS fields (2)..., Our theoretical predictions are compared with data from two public catalogues: the FIREWORKS data for the GOODS-CDF \citep{Wuyts2008} and the NEWFIRM Medium Band Survey data for the COSMOS and AEGIS fields \citep{Whitaker2011}.232 The wide photometric coverage of hese two datasets results in robust ancl relatively precise ohotometric redshift measurements., The wide photometric coverage of these two datasets results in robust and relatively precise photometric redshift measurements.233 Our theoretical samples are selected: using photometric criteria very similar. to hose defining the observed. samples. although we plot the distribution of their true redshifts rather than attempting to reproduce the observational redshift estimation procedure.," Our theoretical samples are selected using photometric criteria very similar to those defining the observed samples, although we plot the distribution of their true redshifts rather than attempting to reproduce the observational redshift estimation procedure."234 Selection by observed-frame {νο ρα. magnitude picks ealaxies on the basis of their rest-frame A ας emission at low redshift. their rest-frame / band emission at 2~I. and their rest-frame optical emission at z 1.5.," Selection by observed-frame $K_s$ -band magnitude picks galaxies on the basis of their rest-frame $K$ -band emission at low redshift, their rest-frame $J$ band emission at $z\sim 1$, and their rest-frame optical emission at $z>1.5$ ."235" For both magnituce limits. the saniples are dominated by intrinsically aint objects at. [ow-redshift but. by intrinsically bright &ealaxiesa bevond 2~1 (for A,« 21.8) or z—1.5 (for Ay<) 23.3))."," For both magnitude limits, the samples are dominated by intrinsically faint objects at low-redshift but by intrinsically bright galaxies beyond $z\sim1$ (for $K_s<21.8$ ) or $z\sim1.5$ (for $K_s<23.3)$ )."236 For both apparent magnitude limits the predictions or the two population svnthesis models agree. and for the iehter limit they are consistent with the observations out o à redshift of almost two.," For both apparent magnitude limits the predictions for the two population synthesis models agree, and for the brighter limit they are consistent with the observations out to a redshift of almost two."237 Both underprediet the counts at ueher redshift. with the ellect being slightly larger for the ? model.," Both underpredict the counts at higher redshift, with the effect being slightly larger for the \citet{Bruzual2003} model."238 This may rellect an underabuncdance of intrinsically xight objects (in the optical) in the model at these redshifts. rut could also be due to magnitude and photometric redshift errors in the data which primarily allect the tails of the distribution.," This may reflect an underabundance of intrinsically bright objects (in the optical) in the model at these redshifts, but could also be due to magnitude and photometric redshift errors in the data which primarily affect the tails of the distribution."239 For the fainter apparent. magnitude limit the nmocel clearly overestimates the number of objects over the redshift range 1«z2.5., For the fainter apparent magnitude limit the model clearly overestimates the number of objects over the redshift range $1<z<2.5$.240 These galaxies typically have stellar masses of order a few. LOMAL. and this discrepancy rellects. the overabundance of objects of this mass. and redshift unity that was Lageed by ?.., These galaxies typically have stellar masses of order a few $10^{10}M_\odot$ and this discrepancy reflects the overabundance of objects of this mass and redshift unity that was flagged by \citet{Guo2011}.241 Phe same problem. was identified in earlier versions of the model by 7. ancl ?.., The same problem was identified in earlier versions of the model by \citet{Kitzbichler2007} and \citet{Torre2011}.242 At higher ancl lower redshifts the abundances agree quite well in model ancl data. rellecting the fact that the semi-analytic model wastuned to fit galaxy abundances at low redshift. ancl predicts an abundance of high-mass galaxies which fits observed estimates quite well at high redshift.," At higher and lower redshifts the abundances agree quite well in model and data, reflecting the fact that the semi-analytic model wastuned to fit galaxy abundances at low redshift, and predicts an abundance of high-mass galaxies which fits observed estimates quite well at high redshift."243 For the HUAC 5.8pm selected samples shown in the right panel of Fig. 3.. ," For the IRAC $5.8\rm{\mu m}$ selected samples shown in the right panel of Fig. \ref{fig:zdist}, ,"244there is a significant cilference between the predictions of the two population svnthesis models., there is a significant difference between the predictions of the two population synthesis models.245 While the 2? model predicts a distribution with similar shape to those in the left panel. the 7. model makes a concordant prediction only at z<L2.," While the \citet{Bruzual2003} model predicts a distribution with similar shape to those in the left panel, the \citet{Maraston2005} model makes a concordant prediction only at $z<1.2$."246" Bewond this point there is a ""bump"" and at higher recshift it. predicts roughly 3 times as many galaxies as the ? model.", Beyond this point there is a “bump” and at higher redshift it predicts roughly 3 times as many galaxies as the \citet{Bruzual2003} model.247 A corresponding bump is not. present. in the observational cata which are better described bv the ο model. at least out to 2~2.5.," A corresponding bump is not present in the observational data which are better described by the \citet{Bruzual2003} model, at least out to $z\sim 2.5$."248 The bump in the ? model is caused by strong rest-frame JLfdy emission [from TP-ACD stars associated. with intermediate age stellar populations., The bump in the \citet{Maraston2005} model is caused by strong rest-frame $JHK$ emission from TP-AGB stars associated with intermediate age stellar populations.249 While this emission. brings the predieted numbers of galaxies into rough agreement with the data at the highest redshifts. it results in an overabundance abono ~2 where the observational datasets appear mos robust.," While this emission brings the predicted numbers of galaxies into rough agreement with the data at the highest redshifts, it results in an overabundance at $z\sim 2$ where the observational datasets appear most robust."250 Since this ellect is also present for data in other wavebands. for which PP-AGB emission is not an issue. it suggests that it results from. the semi-analytic mode overpredieting the abundance of the relevant moderate mass galaxies by a substantial factor at this redshift.," Since this effect is also present for data in other wavebands, for which TP-AGB emission is not an issue, it suggests that it results from the semi-analytic model overpredicting the abundance of the relevant moderate mass galaxies by a substantial factor at this redshift."251 In view of this. the agreement achieved. by the ? mode is probably coincidental. resulting from. the overestimatec abundance of moderate mass galaxies being compcrisatlcc by an overestimate of their rest-frame 44A. mass-to-lieh radios.," In view of this, the agreement achieved by the \citet{Bruzual2003} model is probably coincidental, resulting from the overestimated abundance of moderate mass galaxies being compensated by an overestimate of their rest-frame $JHK$ mass-to-light ratios."252 test-Lrame Luminosity. functions. anc colour. distributions as a function. of redshift’ provide direct. estimates of the abundance evolution. of various galaxy types (eg. star-orming/passive. high/low mass).," Rest-frame luminosity functions and colour distributions as a function of redshift provide direct estimates of the abundance evolution of various galaxy types (e.g. star-forming/passive, high/low mass)."253 However. they require accurate redshifts and appropriate photometry if they are to xe determined reliably from observed-Lrame Luxes.," However, they require accurate redshifts and appropriate photometry if they are to be determined reliably from observed-frame fluxes."254 The wide wavelength: coverage of modern surveys. produces. robust shotometric redshifts. and. in addition. allows rest-frame optical and near-infrared) magnitudes to be determined. by interpolation over the full range O<z4. rather than requiringe an uncertain extrapolation based on an SED fit.," The wide wavelength coverage of modern surveys produces robust photometric redshifts, and, in addition, allows rest-frame optical and near-infrared magnitudes to be determined by interpolation over the full range $0<z<4$, rather than requiring an uncertain extrapolation based on an SED fit."255 ? showed that their semi-analytic model reproduces observed z0.1 luminosity Functions in the SDSS g.r./ and > bands.," \citet{Guo2011} showed that their semi-analytic model reproduces observed $z\sim 0.1$ luminosity functions in the SDSS $g, r, i$ and $z$ bands."256 At higher redshift they implemented the recshilt-dependent dust. model of 7.., At higher redshift they implemented the redshift-dependent dust model of \citet{Kitzbichler2007}.257 This reproduces the observed abundance of colour-selected galaxies al 2~ and 2~3 [or the previous version of the semi-analytic model (see ?).., This reproduces the observed abundance of colour-selected galaxies at $z\sim 2$ and $z\sim 3$ for the previous version of the semi-analytic model \citep[see][]{Guo2009}. .258 In Fig., In Fig.259 4 we show the evolution of the D-band luminosity function from z= Oto z—3 forour current semi-analvtic model., \ref{fig:LfB} we show the evolution of the $B$ -band luminosity function from $z=0$ to $z=3$ forour current semi-analytic model.260 Solid. red and. dashed.blue lines. represent. versions withthe ? and 7. stellar population models respectively.," Solid red and dashedblue lines, represent versions withthe \citet{Maraston2005} and \citet{Bruzual2003} stellar population models respectively."261 Fillecl regions give the expected. 10 field-to-lield scatter for surveysof area 1.4 deg?. except that smaller fields (with area 150 arcmin?) were assumed for the 2.5<z3.5 panel. for the L3<2.2.0 panel fainter than 21.0 and for the OS<2κ12 panel fainter than 19.0.," Filled regions give the expected $\sigma$ field-to-field scatter for surveysof area 1.4 $^2$ , except that smaller fields (with area 150 $^2$ ) were assumed for the $2.5<z<3.5$ panel, for the $1.3<z<2.0$ panel fainter than $-21.0$ and for the $0.8<z<1.2$ panel fainter than $-19.0$ ."262of fundamental atinosphierie parameters of individual stars.,of fundamental atmospheric parameters of individual stars.263 In order to help iu such task we present iu this paper a new code:Widths., In order to help in such task we present in this paper a new code:.264 This software. written in PERL. can be used to compute the oof absorption lines as well as to determine continu. »outs. being very helpful to perform stellar population svuthesis following. for example the method developed wBica(1055) and Schinittetal.(1996).," This software, written in PERL, can be used to compute the of absorption lines as well as to determine continuum points, being very helpful to perform stellar population synthesis following, for example the method developed by\citet{bica88} and \citet{schmitt96}."265 This paper is structured as follows: In Sec., This paper is structured as follows: In Sec.266 ?? we escribe the svsteni requirements as well as the minierical procedures behind the code., \ref{code} we describe the system requirements as well as the numerical procedures behind the code.267 The iuput xuwanmeters aud the outputs of the code are discussed iu Sec. ??.., The input parameters and the outputs of the code are discussed in Sec. \ref{run}.268 A comparison with the measures made whit aand “hand-made” measurements are presented in Sec. ??..," A comparison with the measures made whit and “hand-made"" measurements are presented in Sec. \ref{test}."269 The final remarks are mace in Sec. 7?.., The final remarks are made in Sec. \ref{final}.270 The idea behind is to reproduce the “manual” xocedure used to measure the oof absorption lines in a spectrum. as well as to measure nean coutinuunm fluxes in defined regious and compute he σολπα value at line center.," The idea behind is to reproduce the “manual"" procedure used to measure the of absorption lines in a spectrum, as well as to measure mean continuum fluxes in defined regions and compute the continuum value at line center."271" In addition. using he same inputs. it docs exactly reproduce the measured values being “user” incdependeut (e.g. the uncertainties introduced by the user in ""haud operated” procedures are removed). and thus allowing for a better comparison between nuneasured by different users."," In addition, using the same inputs, it does exactly reproduce the measured values being “user"" independent (e.g. the uncertainties introduced by the user in “hand operated"" procedures are removed), and thus allowing for a better comparison between measured by different users."272 was written in Perl allowing auvone to use it without having any problems with software liceuses.," was written in Perl, allowing anyone to use it without having any problems with software licenses."273 It is freely distributed under GNU Ceneral Public ((GLP)., It is freely distributed under GNU General Public (GLP).274 All the libraries used in the code are also free and distributed πιο GLP license., All the libraries used in the code are also free and distributed under GLP license.275 ssource code ca be freely downloaded from beri]abSootH, source code can be freely downloaded from $\sim$ riffel/software.html.276 All requirements to run can easily be installed in any linux machine qun). they are: Iu addition. we call the atteution to the fact that Perlis installed as default in any liuux flavor aud οσα easily be converted to run wader Microsoft Windows system (without plots).," All requirements to run can easily be installed in any linux machine ), they are: In addition, we call the attention to the fact that Perl is installed as default in any linux flavor and can easily be converted to run under Microsoft Windows system (without plots)."277" Iu general. absorption feature indices are composed. by measurements of relative flux in a central waveleneth interval corresponding to the absorption feature cousidercdl (line Imits A; and A,, iu Fig. 2))"," In general, absorption feature indices are composed by measurements of relative flux in a central wavelength interval corresponding to the absorption feature considered (line limits $\lambda_l$ and $\lambda_u$ in Fig. \ref{ew}) )"278 aud two coutimm sidebands passband regions (spectrum ranges - dots ο e ge , and two continuum sidebands passband regions (spectrum ranges - dots - in the boxes of Fig. \ref{ew}) ).279sidebands provide a reference level Cpseudo-contiuuumu. solid line Fig. 2))," Such sidebands provide a reference level (pseudo-continuum, solid line Fig. \ref{ew}) )"280 your Which the streneth of the absorption feature is evaluated (seeWorthevetal.1991.fordetails)., from which the strength of the absorption feature is evaluated \citep[see][for details]{worthey94}.281 computes the pseudo-contimuum. F(A). iu three wavs: (di) a linear reeression ds computed using all he contin points iu the passband regious aud a straight line (v=ax|b) is computed using the regression cocfiicicnts: (d) a straight line is drawn connecting he nüd-poiuts of the flaukiug passband continui reelons: (di) as a cubic spline®..," computes the pseudo-continuum, $F_c(\lambda)$, in three ways: (i) a linear regression is computed using all the continuum points in the passband regions and a straight line (y=ax+b) is computed using the regression coefficients; (ii) a straight line is drawn connecting the mid-points of the flanking passband continuum regions; (iii) as a cubic ."282 The form in which, The form in which283"As it is well known. the clectromagnuetic torque implics an evolution of the pulsar rotation frequency elven by the relation The electromagnetic iudex is 5,4,=3 for pure dipolar magnetic radiation (Ostriker & Camu 1969)). (Manchester & Tavlor 1977)). 1n which case Tn this equation D is the streneth of the iagnetic field. AF is the neutron star radius. à is the angle )etwoeenu the rotation aud the magnetic axes. aud J is the nomenti of inertia with respect to the rotation axis.","As it is well known, the electromagnetic torque implies an evolution of the pulsar rotation frequency given by the relation The electromagnetic index is $n_{em}=3$ for pure dipolar magnetic radiation (Ostriker $\&$ Gunn \cite{ostri}) ), (Manchester $\&$ Taylor \cite{manch}) ), in which case In this equation $ B$ is the strength of the magnetic field, $ R$ is the neutron star radius, $ \alpha $ is the angle between the rotation and the magnetic axes, and $ I$ is the momentum of inertia with respect to the rotation axis."284 An electromagnetic braking iudex different from 3 can be he consequence of pulsar winds. in which particles having aneular momentum are accelerated away from the pulsar (Coldreich & Julian 1969)). (IXaspi et al. 199 0))," An electromagnetic braking index different from 3 can be the consequence of pulsar winds, in which particles having angular momentum are accelerated away from the pulsar (Goldreich $\&$ Julian \cite{gold}) ), (Kaspi et al. \cite{kaspi}) )"285 or cau be xoduced if non dipolar coumponcuts of the magnetic field are present. as a consequence of strong magnetospheric currents (Blandford & Romani 1958)).," or can be produced if non dipolar components of the magnetic field are present, as a consequence of strong magnetospheric currents (Blandford $\&$ Romani \cite{roma}) )."286 Another possibility is that the magnetic moment. aud then the orque applied o the star. varies in time (Blandford & Romani 19885)). (Cheng 1501).," Another possibility is that the magnetic moment, and then the torque applied to the star, varies in time (Blandford $\&$ Romani \cite{roma}) ), (Cheng \cite{chen}) )."287 Muslinov & Page (1996)) consider the ine evolution of the surface magnetic field of the star. which. according to their model. is very low (10°:109 Gy or a newborn neutron star aud then increases due to the olunic diffusion of the Παν trapped iuuer magnetic field.," Muslimov $\&$ Page \cite{musli}) ) consider the time evolution of the surface magnetic field of the star, which, according to their model, is very low $ 10^8\div 10^9~G$ ) for a newborn neutron star and then increases due to the ohmic diffusion of the initially trapped inner magnetic field."288" The resulting electromagnetic braking mdoex is even ἂν B.,,y is the value of the surface maguetic field at he magnetic pole.", The resulting electromagnetic braking index is given by where $ B_{surf}$ is the value of the surface magnetic field at the magnetic pole.289" Clearly. if D; 270.pou, <3."," Clearly, if $ \dot{B}_{surf}>0$, $ n_{em}<3$."290Other nodels based on the secular variation of the magnetic field have been developed. for instance. by Blauclforcd (1983)) and by Camilo (1996)).," Other models based on the secular variation of the magnetic field have been developed, for instance, by Blandford \cite{bland}) ) and by Camilo \cite{cami}) )."291 A different kiud of mocdols is been proposed by Allen & Torvath (1997)) aud by Link & Epstein (1997)). based ou the growth of the augle a between the maguetic moment aud the rotation axis of he star.," A different kind of models has been proposed by Allen $\&$ Horvath \cite{horv1}) ) and by Link $\&$ Epstein \cite{link}) ), based on the growth of the angle $\alpha $ between the magnetic moment and the rotation axis of the star."292 In this case the electromagnetic braking iudex is Tn all these cases the general expression for the observed electromagnetic braking index is simply which is easily obtained deriving Eq.(03)) with respect to time., In this case the electromagnetic braking index is In all these cases the general expression for the observed electromagnetic braking index is simply which is easily obtained deriving \ref{ompem}) ) with respect to time.293" From Eq.(1)) we fine Assundug n24,=const. from Eq.(7)) we find where νου and Qy are the electromagnetic torque function aud the angular velocity at time f£.=0. Ίο now."," From \ref{kem}) ) we find Assuming $n_{em}=const$, from \ref{dk}) ) we find where $K_{em,0}$ and $\Omega_0$ are the electromagnetic torque function and the angular velocity at time $t=0$, i.e. now."294 Then. the rate of variation of the angular velocity can he written as where [νο-=ooLoan ix: a coustant. which is formally equivalent to Eq.(3)).," Then, the rate of variation of the angular velocity can be written as where $K_{em,eff}={K_{em,0}\over{\Omega^{n_{em}-3}_0}}$ is a constant, which is formally equivalent to \ref{ompem}) )."295 Tuteerating Eq.(3)) we obtain the pulsar “characteristic” time. due to the electromagnetic braking which is: beige ©; the initial aneular velocity of the pulsu. Eq.(101))," Integrating \ref{ompem}) ) we obtain the pulsar “characteristic” time, due to the electromagnetic braking which is: being $\Omega_i$ the initial angular velocity of the pulsar. \ref{tem}) )"296 would eive the true age if only the electromagnetic enüssion was responsible for the pulsar ΜΗ, would give the true age if only the electromagnetic emission was responsible for the pulsar spin-down.297" From Eqs.(3..1)). the streugth of the magnetic field. for n4,=3 ds where α ds the angle between the magnetic aud the rotation axis. P,iy is the rate of variation of the pulsar period due to the electromagnetic torque. and we have assunued f=107kgi."," From \ref{ompem}, \ref{kem}) ), the strength of the magnetic field, for $n_{em}=3$, is where $ \alpha $ is the angle between the magnetic and the rotation axis, $\dot{P}_{em}$ is the rate of variation of the pulsar period due to the electromagnetic torque, and we have assumed $I=10^{38}~kg~ m^2$."298 Iu this section we show that the observed braking mdexes can result from the combined action. iun a pulsar. of the magnetic and eravitational torques.," In this section we show that the observed braking indexes can result from the combined action, in a pulsar, of the magnetic and gravitational torques."299 We assume that the chussion of electromagnetic and gravitational radiation are the only iechanisius acting iu the pulsar., We assume that the emission of electromagnetic and gravitational radiation are the only mechanisms acting in the pulsar.300" The observed spiu-dowu rate ean be expressed as where 0,,, is eiven by Eq.(3)). or Eq.(10)). while where e is the ellipticitv of the pulsar."," The observed spin-down rate can be expressed as where $\dot{\Omega}_{em}$ is given by \ref{ompem}) ), or \ref{nomem}) ), while where $\epsilon $ is the ellipticity of the pulsar."301" We note that. in analogy with Eq.(11)). he pulsar ""characteristic age.o due to the gravitational braking. is Iuu the follfollowing.n we will1 considerler the torque functfuuctious IU and Nyy coustaut iu tie."," We note that, in analogy with \ref{tem}) ), the pulsar “characteristic” age, due to the gravitational braking, is In the following, we will consider the torque functions $K_{em}$ and $K_{gw}$ constant in time."302QSO environmental richness because QSOs are rarely found directly in the centres of over-densities.,QSO environmental richness because QSOs are rarely found directly in the centres of over-densities.303 At lower AGN Iuminosities optically selected: Sevífert galaxies seem much less likely to be found in rich environments., At lower AGN luminosities optically selected Seyfert galaxies seem much less likely to be found in rich environments.304 7. analyse a sample of Sevfert galaxies and find no significant dillerence. between the environmental richness. or the probability. of finding a close companion galaxy. Compared to a matched sample of non-active ealaxies.," \citet{1998ApJ...496...93D} analyse a sample of Seyfert galaxies and find no significant difference between the environmental richness, or the probability of finding a close companion galaxy, compared to a matched sample of non-active galaxies."305 Although they do find a dillerence between the environments of the Sevíert 1: ancl Sevfert 2: sub-samples with Sevfert Is being in poorer environments. an observation that they cannot explain in terms of the Unified. Model of AGN. which predicts that there should be no dillerence in the environment of these two AGN classes.," Although they do find a difference between the environments of the Seyfert 1 and Seyfert 2 sub-samples with Seyfert 1s being in poorer environments, an observation that they cannot explain in terms of the Unified Model of AGN, which predicts that there should be no difference in the environment of these two AGN classes."306 In all the above cases the sample sizes have been necessarily small (tvpically several tens of QSOs) because of the limitations in performing large numbers of pointed observations. especially if the sample is at high redshift (see table Lol?) for à summary of a sample of studies of AGN environments).," In all the above cases the sample sizes have been necessarily small (typically several tens of QSOs) because of the limitations in performing large numbers of pointed observations, especially if the sample is at high redshift (see table 1 of \citet{2001AJ....122...26B}307 for a summary of a sample of studies of AGN environments)."308 The situation at low redshifts (2< O.1). however. is now somewhat alleviated by large spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS) or 2df£ Galaxy Redshift Survey. which include many thousands of AGN.," The situation at low redshifts $z<0.1$ ), however, is now somewhat alleviated by large spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS) or 2df Galaxy Redshift Survey, which include many thousands of AGN."309 Despite this plethora of cata different stuclics still disagree to some extent on some details of the environments of AGN., Despite this plethora of data different studies still disagree to some extent on some details of the environments of AGN.310 ? find essentially no change in the fraction of galaxies with an AGN. across nearly two decades in environmental density.," \citet{2003ApJ...597..142M} find essentially no change in the fraction of galaxies with an AGN, across nearly two decades in environmental density."311 Of. the nearly 5000. ealaxics studied up to —40percent showed some sign of nuclear activity (an upper limit based on modelling of the lower S/N emission lines: the higher S/N lines allowed ~20percent to be unambiguously classified. as AGN). the fraction remaining constant with densitv.," Of the nearly 5000 galaxies studied up to $\sim40~per~cent$ showed some sign of nuclear activity (an upper limit based on modelling of the lower S/N emission lines; the higher S/N lines allowed $\sim20~per~cent$ to be unambiguously classified as AGN), the fraction remaining constant with density."312 On the other hand. star-forming galaxies are found in much greater abundance in rarefied environments — the so called SER-clensity relation.," On the other hand, star-forming galaxies are found in much greater abundance in rarefied environments – the so called SFR-density relation."313 Passive galaxies. of course. are Found in greater abundances in denser environments.," Passive galaxies, of course, are found in greater abundances in denser environments."314 Such a high. and constant. fraction of galaxies containing an AGN rather suggests that the fuelling mechanism for these lower Iuminosity objects (mostly LINERS. the most common and lowest) luminosity ACN elass) is à frequent occurrence. and common to a large range of environments.," Such a high, and constant, fraction of galaxies containing an AGN rather suggests that the fuelling mechanism for these lower luminosity objects (mostly LINERS, the most common and lowest luminosity AGN class) is a frequent occurrence, and common to a large range of environments."315 Major mergers therefore seem highly unlikely as à common fuelling mechanism. as do any other cluster related mechanis," Major mergers therefore seem highly unlikely as a common fuelling mechanism, as do any other cluster related mechanisms."316nis. ? also use the SDSS data to study the AGN fraction as a Function of environmental density., \citet{2004MNRAS.353..713K} also use the SDSS data to study the AGN fraction as a function of environmental density.317 Phey linc a somewhat different result. from ? in that twice as many galaxies vost AGN in low density environments as in hieh. a trend hey attribute to the fact that ACN and. star-formation are related in some way.," They find a somewhat different result from \citet{2003ApJ...597..142M} in that twice as many galaxies host AGN in low density environments as in high, a trend they attribute to the fact that AGN and star-formation are related in some way."318 However. their. classification of AGN clillers from that of 2.. which is possibly the cause of he difference.," However, their classification of AGN differs from that of \citet{2003ApJ...597..142M}, which is possibly the cause of the difference."3192? only study AGN with OHI] luminosities lo’L. (total fraction  (0.1). whereas ? study AGN with a wider range of OHI] luminosities. resulting in the jigher overall AGN fraction.,"\citet{2004MNRAS.353..713K} only study AGN with O[III] luminosities $>10^{7}~L_{\odot}$ (total fraction $\sim0.1$ ), whereas \citet{2003ApJ...597..142M} study AGN with a wider range of O[III] luminosities, resulting in the higher overall AGN fraction."320 Whereas AGN with high OLL] unminosities show an environmental dependence. those with ower OLI] luminosities do not. so in-fact the two results are not in contradiction.," Whereas AGN with high O[III] luminosities show an environmental dependence, those with lower O[III] luminosities do not, so in-fact the two results are not in contradiction."321 7 extend the work of ? to à larger sample size and find much the same AGN fraction (18per cent)., \citet{2004ApJ...610L..85W} extend the work of \citet{2003ApJ...597..142M} to a larger sample size and find much the same AGN fraction $18~per~cent$ ).322 They study the two-point correlation Function of AGN and compare it to all ealaxies. finding no significant dillerence between the two. sugeesting that AGN follow the distribution of the normal galaxy population. and are thus unbiased tracers of mass in the Universe.," They study the two-point correlation function of AGN and compare it to all galaxies, finding no significant difference between the two, suggesting that AGN follow the distribution of the normal galaxy population, and are thus unbiased tracers of mass in the Universe."323 A-rayv surveys are very observationally ellicicnt at. [finding AGN over à wide range in redshift., X-ray surveys are very observationally efficient at finding AGN over a wide range in redshift.324 Hard X-ray luminosity. in particular is a highly unbiased measure of AGN power. as the only thing that is being probed is the accretion rate of the SBIL itself: the details of the exact AGN type and viewing angle are unimportant due to the penetrating power of hard N-ravs.," Hard X-ray luminosity in particular is a highly unbiased measure of AGN power, as the only thing that is being probed is the accretion rate of the SBH itself; the details of the exact AGN type and viewing angle are unimportant due to the penetrating power of hard X-rays."325 The narrow ΟΠΗ) emission line is also thought to be an unbiased. tracer of ACGIN. because it originates from bevond the obscuring torus. and. is commonly used as a measure of activity in low-z ACN.," The narrow [OIII] emission line is also thought to be an unbiased tracer of AGN because it originates from beyond the obscuring torus, and is commonly used as a measure of activity in low-z AGN."326 The fairly tight correlation between hare X-ray anc OLLI] luminosity (2) indicates that the same physical process is likely to be responsible for both. namely the accreting SBI.," The fairly tight correlation between hard X-ray and [OIII] luminosity \citep{1999AJ....118.1169X} indicates that the same physical process is likely to be responsible for both, namely the accreting SBH."327 However. at higher redshifts the ΟΠΠ line becomes harder to detect in weak ACN as more of the galaxy. light falls into the slit or fibre aperture. washing out the nuclear light.," However, at higher redshifts the [OIII] line becomes harder to detect in weak AGN as more of the galaxy light falls into the slit or fibre aperture, washing out the nuclear light."328 This problem does not allect the X-ray emission from AGN. however.," This problem does not affect the X-ray emission from AGN, however."329 ανά X-ray emission is an excellent tracer of AGN activity because it is dillicult for anvthing other than a SBIL to generate. X-ray. luminosities in excess of 1077. ere st (210 keV)., Hard X-ray emission is an excellent tracer of AGN activity because it is difficult for anything other than a SBH to generate X-ray luminosities in excess of $10^{42}$ erg $^{-1}$ $2-10$ keV).330 Hare X-rays are also allected far less w intrinsic absorption than soft. N-rays and can penetrate aree column densities of intervening neutral hydrogen (up o 7107 em.2 7) that would. essentially completely absorb οποίος of energy less than 2 keV. Of course. nothing is »erfect and for extremely high column densities. resulting in Compton thick obscuration (Ng~1.5107 7). even ward X-rays are absorbed.," Hard X-rays are also affected far less by intrinsic absorption than soft X-rays and can penetrate large column densities of intervening neutral hydrogen (up to $\sim 10^{23}$ $^{-2}$ ) that would essentially completely absorb photons of energy less than $2$ keV. Of course, nothing is perfect and for extremely high column densities, resulting in Compton thick obscuration $N_{H}\sim 1.5\times33110^{24}$ $^{-2}$ ), even hard X-rays are absorbed."332 But for the purposes of this study we shall ignore Compton thick GN. with the assumption hat they constitute a relatively small fraction of the total population (see 7. for a discussion of the Compton thick contribution)," But for the purposes of this study we shall ignore Compton thick AGN, with the assumption that they constitute a relatively small fraction of the total population (see \citet{2003ApJ...598..886U} for a discussion of the Compton thick contribution)."333 The details of the data. its reduction and the source detection algorithmapplied. to it are. described. in. both papers 1 and 2.," The details of the data, its reduction and the source detection algorithmapplied to it are described in both papers 1 and 2."334 Briellv. the exposures are of ~50 ks duration. reaching a210 keV N-ray [lux completeness limit of ~61015 ere ocn 2404c.," Briefly, the exposures are of $\sim 50$ ks duration, reaching a $2-10$ keV X-ray flux completeness limit of $\sim 6 \times33510^{-15}$ erg $^{-2}$ $^{-1}$."336 The X-ray point sources were identified with optical objects from the CEDE. catalogues as described in detail in paper 2., The X-ray point sources were identified with optical objects from the CFDF catalogues as described in detail in paper 2.337 Photometric redshifts were determined. from the multi-bancd photometry ofthe reliably identified: AGN., Photometric redshifts were determined from the multi-band photometry ofthe reliably identified AGN.338 These AGN form the basis for this present work., These AGN form the basis for this present work.339 The CEDE catalogues were derived. from a campaign to image one square degree to Zig(39.37)~\25.5. with comparable depths in C. D and V.," The CFDF catalogues were derived from a campaign to image one square degree to $I_{AB}(3\sigma,3\arcsec)\sim25.5$, with comparable depths in $U$ , $B$ and $V$ ."340 Ehe survey area was split into 4 sub-surveys of 30 307. two of which were used for," The survey area was split into 4 sub-surveys of $30\arcmin \times 30\arcmin$ , two of which were used for"341of 80420 M. yr!.,of $\pm$ 20 $M_\odot$ $^{-1}$.342 Thus. to within a factor of ~2 and within the uncertainties. the two estimates are consistent.," Thus, to within a factor of $\sim$ 2 and within the uncertainties, the two estimates are consistent."343 The optically derived SFR. on the other hand. is significantly lower than that predicted from the total X-ray luminosity.," The optically derived SFR, on the other hand, is significantly lower than that predicted from the total X-ray luminosity."344 From the calibration of Ranalli et al. (, From the calibration of Ranalli et al. (3452003). 10 keV luminosity of 3x107 erg s7!. if entirely attributed to star formation. requires a SFR of =300 M. !. a factor of ~5 higher than our estimates based on optical line emission.,"2003), a 0.5--10 keV luminosity of $\times$ $^{42}$ erg $^{-1}$, if entirely attributed to star formation, requires a SFR of $\approx$ 300 $M_\odot$ $^{-1}$, a factor of $\sim$ 5 higher than our estimates based on optical line emission."346 There are two possible explanations for this apparent discrepancy., There are two possible explanations for this apparent discrepancy.347" First. the optically derived SFR may be an underestimate of the total SER if star formation is more extended than the 2""-wide slit: indeed. in case the distribution of star-forming regions is wider than that of the optical continuum emission. the correction for slit losses underestimates the actual total flux of emission lines."," First, the optically derived SFR may be an underestimate of the total SFR if star formation is more extended than the $''$ -wide slit: indeed, in case the distribution of star-forming regions is wider than that of the optical continuum emission, the correction for slit losses underestimates the actual total flux of emission lines."348 By contrast. the X-ray emission integrates over a larger area.," By contrast, the X-ray emission integrates over a larger area."349 Second. only a portion (~20%)) of the X- emission originates from star formation. with the dominant fraction coming from an AGN.," Second, only a portion $\sim$ ) of the X-ray emission originates from star formation, with the dominant fraction coming from an AGN."350 Although the optical spectrum shows no signs of nonstellar activity. we cannot exclude the presence of a low-level AGN. especially one that is heavily obscured.," Although the optical spectrum shows no signs of nonstellar activity, we cannot exclude the presence of a low-level AGN, especially one that is heavily obscured."351 Both of these possibilities can be tested with further observations., Both of these possibilities can be tested with further observations.352 The detection of |O n]. [O mir]. and Ημ also allows us to infer the gas-phase oxygen abundance.," The detection of [O ], [O ], and $_\beta$ also allows us to infer the gas-phase oxygen abundance."353 Following Kobulnicky et al. (, Following Kobulnicky et al. (3541999). the Ro: parameter gives 12 + log (O/H) = 7.6 or 8.9.,"1999), the $R_{\rm 23}$ parameter gives 12 + log (O/H) = 7.6 or 8.9."355 Considering the intrinsic luminosity of the source (rest-frame Myx —21.5 mag) and its |O m]/[N i] ratio. the larger of the two values is likely to be the correct one (see Kobulnicky et al.," Considering the intrinsic luminosity of the source (rest-frame $M_B356\approx$ $-$ 21.5 mag) and its [O ]/[N ] ratio, the larger of the two values is likely to be the correct one (see Kobulnicky et al."357 1999). pointing to a basically solar oxygen abundance.," 1999), pointing to a basically solar oxygen abundance."358 The comparison of the USNO archival magnitudes. with the results of the optical photometry acquired in Loiano seems to suggest long-term variability for the source., The comparison of the USNO archival magnitudes with the results of the optical photometry acquired in Loiano seems to suggest long-term variability for the source.359 However. we caution the reader that it is not infrequent to find discrepancies as large as 0.6 mag (see e.g. Masetti et al.," However, we caution the reader that it is not infrequent to find discrepancies as large as 0.6 mag (see e.g. Masetti et al."360 2003) between the automatically extracted archival magnitudes and the actual ones. calibrated through Landolt (1992) fields.," 2003) between the automatically extracted archival magnitudes and the actual ones, calibrated through Landolt (1992) fields."361 This 1s most likely due to the fact that magnitude extraction from archival images was performed by using automatic pipelines and was (for the oldest surveys) applied to photographic plates., This is most likely due to the fact that magnitude extraction from archival images was performed by using automatic pipelines and was (for the oldest surveys) applied to photographic plates.362 In the minisurvey in Foschini et al. (, In the minisurvey in Foschini et al. (3632002b) 18 ULXs in 10 nearby galaxies were detected.,2002b) 18 ULXs in 10 nearby galaxies were detected.364 According to the statistics from the Lockman hole studies (Hasinger et al., According to the statistics from the Lockman hole studies (Hasinger et al.365 2001) we expect to find 3.2 background sources with 2-10 keV flux higher than 107 ere em s! in an area corresponding to the sum of all the Das ellipses of the considered galaxies., 2001) we expect to find 3.2 background sources with 2–10 keV flux higher than $10^{-14}$ erg $^{-2}$ $^{-1}$ in an area corresponding to the sum of all the $D_{25}$ ellipses of the considered galaxies.366 If we also take into account the case reported in Foschini et al. (, If we also take into account the case reported in Foschini et al. (36720022). with the present identification we already founc 2 background objects among the 18 ULX candidates: so. we are close to the statistical limit by assuming that the Lockman Hole studies are valid over a wide range of the sky.,"2002a), with the present identification we already found 2 background objects among the 18 ULX candidates: so, we are close to the statistical limit by assuming that the Lockman Hole studies are valid over a wide range of the sky."368 Albeit the two selected cases might suffer from an observational bias (they are quite bright and outside the mam body of the parent galaxy). additional detection of other background objects in the minisurvey should be considered with care.," Albeit the two selected cases might suffer from an observational bias (they are quite bright and outside the main body of the parent galaxy), additional detection of other background objects in the minisurvey should be considered with care."369 In conclusion. we wish to remark that this ease. as that of the ULX in NGC 4698 reported by Foschini et al. (," In conclusion, we wish to remark that this case, as that of the ULX in NGC 4698 reported by Foschini et al. ("3702002a). has been instructive: despite the availability of broad-band multiwavelength data. they have proven insufficient to disclose the true nature of the source.,"2002a), has been instructive: despite the availability of broad-band multiwavelength data, they have proven insufficient to disclose the true nature of the source."371 The most crucial piece of information is optical spectroscopy of sufficient quality to reveal clear spectral features. which simultaneously provide redshift information and spectral diagnostics.," The most crucial piece of information is optical spectroscopy of sufficient quality to reveal clear spectral features, which simultaneously provide redshift information and spectral diagnostics."372"We consider the possibility of a combination of stellar insolation and tidal “boosting” raising the surface temperature of a moon into the temperate range (273<T,,373 IK).",We consider the possibility of a combination of stellar insolation and tidal “boosting” raising the surface temperature of a moon into the temperate range $273 < T_{eq} <373$ K).373 A very general set of constraints is presented: demonstrating how this can lead {ο an extended temperate. or habitable. zone of as uch as a factor ~2 greater distance Irom (he parent star lor massive (~LAL ) moons.," A very general set of constraints is presented; demonstrating how this can lead to an extended temperate, or habitable, zone of as much as a factor $\sim 2$ greater distance from the parent star for massive $\sim 1 $ $_{\oplus}$ ) moons."374 We also find that the relevant orbital terrain around the known exoplanets [ον tidallv boosted. temperate. massive moons is essentially tlie same as that of the Galilean satellites in (he Jovian svstem.," We also find that the relevant orbital terrain around the known exoplanets for tidally boosted, temperate, massive moons is essentially the same as that of the Galilean satellites in the Jovian system."375 However. the required. tidal heating energv budgets range from the level seen in Io to as much as 100 times greater.," However, the required tidal heating energy budgets range from the level seen in Io to as much as 100 times greater."376 At this upper ex(reme il is not clear whether such dissipation could be either loue-lived (more than a few 100 million years) or compatible with a habitable surface environment., At this upper extreme it is not clear whether such dissipation could be either long-lived (more than a few 100 million years) or compatible with a habitable surface environment.377 We have ignored a multitude of possible additional factors in making our estimates - for example. there are likely significant resonance conditions which max occur for moons orbiüng planets which are themselves in eccentric orbils about the parent star. aud (there are also likely influences [rom other. as vel undetected. planets in the svstem.," We have ignored a multitude of possible additional factors in making our estimates - for example, there are likely significant resonance conditions which may occur for moons orbiting planets which are themselves in eccentric orbits about the parent star, and there are also likely influences from other, as yet undetected, planets in the system."378 We have also nol considered. resonances bevond simple orbital mean-motions - for exemple. librational resonances such as that likely in the case of Enceladus. (Wisdom2004).," We have also not considered resonances beyond simple orbital mean-motions - for example, spin-orbit librational resonances such as that likely in the case of Enceladus \citep{wisdom04}."379. such effects are interesting. and should be explored further., Such effects are interesting and should be explored further.380 We have also not included anv estimate of thekelihood that moons will enter into resonance conditions that will dive orbital eccentricity and hence moon-planet σα] dissipation - or of the dynamical timescales of such situations., We have also not included any estimate of the that moons will enter into resonance conditions that will drive orbital eccentricity and hence moon-planet tidal dissipation - or of the dynamical timescales of such situations.381 This will require a more extensive study which should include modeling of (he origin of a satellite or moon svstem and the potential orbital migration of major satellites within such a svstem., This will require a more extensive study which should include modeling of the origin of a satellite or moon system and the potential orbital migration of major satellites within such a system.382 Such a study should. also include an evaluation of anv variation in the composition and structural properties of the satellites resulting from different. circumstellar/circumplanetary Formation distances., Such a study should also include an evaluation of any variation in the composition and structural properties of the satellites resulting from different circumstellar/circumplanetary formation distances.383 llowever. taken al lace value. one of the implications of the above observations is that the potential for sub-surface oceans in i¢ev moons. or even ocean moons (lor those svstems which lie close enough to the parent star for significant surface temperatures to develop. boosted by Udal dissipation) suggests that. broad questions of habitability may need to be re-visited to include such environments as signilicant potential biospheres.," However, taken at face value, one of the implications of the above observations is that the potential for sub-surface oceans in icy moons, or even ocean moons (for those systems which lie close enough to the parent star for significant surface temperatures to develop, boosted by tidal dissipation) suggests that broad questions of habitability may need to be re-visited to include such environments as significant potential biospheres."384 As the parameter space for exoplanet detection expands to include lower-mass. larger-orbilalradii svstems. and if future missions and instruments begin to detect the presence of moons around giant planets. we will be able to extend (he evaluation of such potential habitats.," As the parameter space for exoplanet detection expands to include lower-mass, larger-orbital-radii systems, and if future missions and instruments begin to detect the presence of moons around giant planets, we will be able to extend the evaluation of such potential habitats."385 The author acknowledges the funding support of the Columbia Astrobiology Center through Columbia Universitys Initiatives in Science ancl Engineering. and (he support οἱ ihe Columbia Astrophysics Laboratory.," The author acknowledges the funding support of the Columbia Astrobiology Center through Columbia University's Initiatives in Science and Engineering, and the support of the Columbia Astrophysics Laboratory."386 This work is also directly supported by a NASA Astrobiology: Exobiologv and Evolutionary Diologv: and Planetary Protection Research, This work is also directly supported by a NASA Astrobiology: Exobiology and Evolutionary Biology; and Planetary Protection Research387 function ΑΙ 0)= sin? 0.,where Each term in \ref{du_dm}) ) is manifestly positive.388 determined b, This completes the proof.389"ytheconstraint F=0. caustics for positive 6 No r labeluo,austic We now prove therather remarkable result that quasi-spherical light cones arefre"," The fact that adding mass to the Kerr field should the focusing of ingoing radial light rays is at first sight paradoxical, but can be explained by the circumstance that the material source is not at the origin, but located on the singular equatorial ring $x^2+y^2=a^2$ , $z=0$ (see \ref{oblate}) ))."390e ofcaustics forall positive valuesof theKerr radial coordinate r. This is trivially tr," The source has a peculiar, “demi-pole” structure \cite{bi}, made possible by the double-sheeted structure of the Kerr manifold: a ring of positive mass in the sheet $r>0$ is bonded to a ring of negative mass in the sheet $r<0$."391ueif i= 0.when these surfaces are siu, Light rays heading inwards toward the origin in the sheet $r>0$ are deflected outward by the ring and defocused.392ply Πο cones inMinkoseski space with vertices atthe spatial origin. repr, Repulsive effects have become dominant by the time they pass through the disk $r=0:$ the rays are then refocused and form a caustic in the sheet $r<0$.393"esented by r= 0.0 =0or zinoblate spheroidal coordinates according to(16)). We shall prove thatit istrue« 201,for"," This description is more than just hand-waving: the Keres Newtonian analogue model of the Kerr field \cite{ki} has a source structure and equations of motion closely resembling Kerr (apart from frame dragging), and displays precisely this behaviour."394tiori forii 20by effectively shiowiug, Figure \ref{fig1} (a result of numerical integrations described in Sec.395 converge less rapidly as As noted inSec.L.. formation o," 10) shows the ingoing generators mapped onto the flat background of the Kerr-Schild decomposition, using the rectangular coordinates defined in \ref{oblate}) ). ("396facaustic alon,Effects of frame dragging are not shown).397gagenerator issignalled by pPOQsu > 0. (10) We DU. x =P(A.0).) VA sin?," According to the Keres-Kerr model, the gravitational “force” should become repulsive for $ 398r \approx a$, and, indeed, the generators have inflection points near this radius."399 Since ϐ decreases with +atfixedA accordingto (25)).the factorP=a relmain positive.Any possiblecaustic inthe Kerr positive-rsheet cannot Thus. noueofthe three factorsp. Qand sindcanreach zero soonerfor positive i) than theydo space.and," Since the quasi-spherical light cones $t\pm\rb=constant$ are characteristics of the wave operator, the coordinates $\rb,\lambda$ are well adapted for representing asymptotically spherical high-frequency The wave equation for $\Psi(\rb,\lambda,\phi,t)$ onthe Kerr background \ref{kerr_zamo}) ) is where $ \gamma=\mu P Q \sin\theta$ gives the area element (degenerate volume element) on the light cone \ref{dsll}) Introducing the ansatz into \ref{wave}) ), we find where has no explicit $\omega$ -dependence."400 hence theydonot according Opi — ——=oOF, We can re-express $\omega_B$ in \ref{wave2}) ) as a partial derivative withrespect to $\rb$ .401 al alI. (15), Defining an angular function402function (CCE) of the time series in different energy ranges.,function (CCF) of the time series in different energy ranges.403 The CCF is defined in most textbooks of statistical analvsis (soeforexample?)..," The CCF is defined in most textbooks of statistical analysis \citep[see for404example][]{Bevington_Robinson_2003}."405 In fig., In fig.406 2. we show the absolute value of the CCF between the MCAL lightcurve in the 0.3 - 1 MeV. enerex baud aud the lishteurves in 1 5 MeV. (top panel) aud above 5 MeV (bottom panel)., \ref{fig:MCAL_cross-cor} we show the absolute value of the CCF between the MCAL lightcurve in the 0.3 - 1 MeV energy band and the lightcurves in 1 – 5 MeV (top panel) and above 5 MeV (bottom panel).407 We use a time lag between—30 8 and |160 S and a bin size of 1.021 s (the same as in fig. 1)).," We use a time lag between$-30$ s and $+160$ s and a bin size of 1.024 s (the same as in fig. \ref{fig:MCAL-GRID_lightcurve}) ),"408 the shortest allowed by the statistics of the data., the shortest allowed by the statistics of the data.409 The peak iu the CCF is taken as the measure of spectral lag (ες , The peak in the CCF is taken as the measure of spectral lag \citep{Norris_2002}. .410We can see from fig., We can see from fig.411 2 that the time lag in the MCAL liehteurves is less thaw INDIES The position of the maxima cussion dn the first Dip does not change significautlv in the three enerev bands. as demonstrated with the cross-correlation function.," \ref{fig:MCAL_cross-cor} that the time lag in the MCAL lightcurves is less than 1.024 s. The position of the maximum emission in the first bump does not change significantly in the three energy bands, as demonstrated with the cross-correlation function."412 Conversely. the first dump width shows variations as a fiction of caerey. measured by assuniug a svuunetric Caussian shape.," Conversely, the first bump width shows variations as a function of energy, measured by assuming a symmetric Gaussian shape."413 Frou the Gaussian fit we found. again. that the position. does not chauge senificautle with energwv. while the width. consistent iu the first and second bands. is smaller in the highest energv— interval.," From the Gaussian fit we found, again, that the position does not change significantly with energy, while the width, consistent in the first and second bands, is smaller in the highest energy interval."414" The Gaussian fiction is just a qualitative model for the bunip shape. aud the high values of the reduced chi square, Which are not formally acceptable. are dominated by the internal peak structure."," The Gaussian function is just a qualitative model for the bump shape, and the high values of the reduced chi square, which are not formally acceptable, are dominated by the internal peak structure."415 Simularly. the position of the masini cussion in the second main biuup docs wot change significantly in the three energy. bauds but. in this case. the Inunp shape is dominated by the superposition of iiv overlapping peaks aud thus cannot be modelled using the same method as applied to the first bunip.," Similarly, the position of the maximum emission in the second main bump does not change significantly in the three energy bands but, in this case, the bump shape is dominated by the superposition of many overlapping peaks and thus cannot be modelled using the same method as applied to the first bump."416 The GRID lishteurve shows two broad bunps. at the sale position as iu MCAL. without evident features in the iuterbunup region (sec fg. 1)).," The GRID lightcurve shows two broad bumps, at the same position as in MCAL, without evident features in the interbump region (see fig. \ref{fig:MCAL-GRID_lightcurve}) )."417 The average background level iu the GRID time series around the time interval of CRB 100721B is 0.12etus|l corresponding to a fux of L«10↽phem22s! when corrected for. the effective⋅⋅ area! at LOO MeV. for an average off-axis angle of 50°.," The average background level in the GRID time series around the time interval of GRB 100724B is $0.12 \; \mathrm{cts \; s^{-1}}$, corresponding to a flux of $4418\times 10^{-4} \; \mathrm{ph \; cm^{-2} \; s^{-1}}$ when corrected for the effective area at 100 MeV for an average off-axis angle of $50 \degree$."419 The ouset of the eiua ταν cussion does not slow any significant delavs with respect to the hard N-ray band: in fact the first photon is detected by GRID at fy|10.9 s. during the rise of the first buinirp in MCAL. and has an energv of z220 MeV. The highest energy photon (23.5 GeV) is detected by GRID inunediatelv before the muaxinmuu of the second bump in the MCÀL helteurve.," The onset of the gamma ray emission does not show any significant delays with respect to the hard X-ray band; in fact the first photon is detected by GRID at $t_0 + 10.9$ s, during the rise of the first bump in MCAL, and has an energy of $\simeq 220$ MeV. The highest energy photon $\simeq 3.5$ GeV) is detected by GRID immediately before the maximum of the second bump in the MCAL lightcurve."420 The narrow peak around &fy|35 s in the MCAL lightcurve above 5 MeV is also evident iu the CRID data., The narrow peak around $\simeq t_0 + 35$ s in the MCAL lightcurve above 5 MeV is also evident in the GRID data.421 The absolute value of the CCF between the time series of the GRID and MCAL (vithout background subtraction) in the three energy bands (0.3 - 1 AIeV. 1 5 MeV. aud above 5 MeV) is shown in fig. 3..," The absolute value of the CCF between the time series of the GRID and MCAL (without background subtraction) in the three energy bands (0.3 - 1 MeV, 1 – 5 MeV, and above 5 MeV) is shown in fig. \ref{fig:MCAL-GRID_cross-cor}."422 In the cross-correlation we used for all the shortestlieliteurves the same time bin of 2.5 s which is the value allowed bv the statistics of the CRID data and the same one as used for the lighteurve in fie. 1..," In the cross-correlation we used for all the lightcurves the same time bin of 2.5 s, which is the shortest value allowed by the statistics of the GRID data and the same one as used for the lightcurve in fig. \ref{fig:MCAL-GRID_lightcurve}. ."423 Since the GRB is within the GRID FoV only from ty|6 suntil ty|125 s. we used," Since the GRB is within the GRID FoV only from $t_0 + 6$ suntil $t_0 + 125$ s, we used"424atmosphere. resulting in a more evenly distributed. chemical composition.,"atmosphere, resulting in a more evenly distributed chemical composition."425 In the upper atmosphere (he ellects of UV photolvsis and molecular diffusion dominate., In the upper atmosphere the effects of UV photolysis and molecular diffusion dominate.426 Chemically. the main source ofdifferences between the predictions of equilibrium chemistry and our photochemical modeling is (hiat certain reactions occur so slowly that the chemistry will never converge to equilibrium before photolvsis ancl müxing take over.," Chemically, the main source of differences between the predictions of equilibrium chemistry and our photochemical modeling is that certain reactions occur so slowly that the chemistry will never converge to equilibrium before photolysis and mixing take over."427 Specilicallv. photolvsis of both anunonia and methane occur readily in (he upper atmosphere. but reactions that create (hese molecules are not [avored.," Specifically, photolysis of both ammonia and methane occur readily in the upper atmosphere, but reactions that create these molecules are not favored."428 The end result is that. nitrogen. and carbon preferentially form into molecules other than methane and ammonia. resulting in lower abundances of both of these species in the upper atnosphere than what is predicted. by equilibrium calculations.," The end result is that nitrogen and carbon preferentially form into molecules other than methane and ammonia, resulting in lower abundances of both of these species in the upper atmosphere than what is predicted by equilibrium calculations."429 Nitrogen tends to combine into the very stable No molecule and also forms a smaller aànount of ICN., Nitrogen tends to combine into the very stable $_2$ molecule and also forms a smaller amount of HCN.430 Carbon is redistributed into a wide varietv of molecules including CO. COs. CsII» (acetvlene). Coll; (ethylene). Coll; (ethane). and ICN (hydrogen evanide). resulüng in a complex carbon chemistry as shown in Figure 3..," Carbon is redistributed into a wide variety of molecules including CO, $_2$ , $_2$ $_2$ (acetylene), $_2$ $_4$ (ethylene), $_2$ $_6$ (ethane), and HCN (hydrogen cyanide), resulting in a complex carbon chemistry as shown in Figure \ref{f3}."431 In our models using the quiet M-star as the UV input spectrum (dashed lines in Figures and 3)) very little UV photolvsis takes place., In our models using the quiet M-star as the UV input spectrum (dashed lines in Figures \ref{f2} and \ref{f3}) ) very little UV photolysis takes place.432 These results can therefore be interpreted as showing the effects of vertical mixing in the absence of any significant. photochemistry., These results can therefore be interpreted as showing the effects of vertical mixing in the absence of any significant photochemistry.433 In (hese models. (he atmosphere is generally well-mixed. and abundance profiles for most major consütuents are constant with height throughout most of the atmosphere until molecular diffusion takes over at high altitudes.," In these models, the atmosphere is generally well-mixed, and abundance profiles for most major constituents are constant with height throughout most of the atmosphere until molecular diffusion takes over at high altitudes."434 This is in sharp contrast to the predictions of equilibrium chemistry that produce strong gradients in (he abundances of several kev molecules includiusg CO. COs. and NIL; (dotted lines in Figures 2. and 3)).," This is in sharp contrast to the predictions of equilibrium chemistry that produce strong gradients in the abundances of several key molecules including CO, $_2$, and $_3$ (dotted lines in Figures \ref{f2} and \ref{f3}) )."435 Our model results produce obvious quenchecl behavior for several molecules at modest depth including CO. COs and ICN.," Our model results produce obvious quenched behavior for several molecules at modest depth including CO, $_2$ and HCN."436 For these molecules. abundance profiles clearly follow their equilibrium values at depth. whereas their abundances ave fairlv constant above the quench point.," For these molecules, abundance profiles clearly follow their equilibrium values at depth, whereas their abundances are fairly constant above the quench point."437 All of these molecules also clisplay a gradual transition between (he quenched regime and the equilibrium regime. which is a result of atmospheric mixing.," All of these molecules also display a gradual transition between the quenched regime and the equilibrium regime, which is a result of atmospheric mixing."438 In the absence of mixing. the abundance profiles would transilion sharply from equilibrium (o their constant. quenchecl abundances at the quench point.," In the absence of mixing, the abundance profiles would transition sharply from equilibrium to their constant quenched abundances at the quench point."439 CO. COs. and IICN display fully quenched behavior above 1-10 bar.," CO, $_2$, and HCN display fully quenched behavior above 1-10 bar."440 Ia contrast. No and NIT; display quenched behavior even al the very. base of the atmosphere for all of our models. which implies that the actual quench point for these species lies deeper than. 1000 bar.," In contrast, $_2$ and $_3$ display quenched behavior even at the very base of the atmosphere for all of our models, which implies that the actual quench point for these species lies deeper than 1000 bar."441 For (his reason. it is possible that the actual abundances of nitrogen-bearing species in GJ 1214b's atmosphere may differ somewhat [rom the values we report here.," For this reason, it is possible that the actual abundances of nitrogen-bearing species in GJ 1214b's atmosphere may differ somewhat from the values we report here."442 Without high levels of UV. irradiation. methane remains the dominant. carbon-bearing species throughout the atmosphere and is present at high abundances ranging from for solar metallicity atmospheres to when the metallicity is enhanced to 30x solar.," Without high levels of UV irradiation, methane remains the dominant carbon-bearing species throughout the atmosphere and is present at high abundances ranging from for solar metallicity atmospheres to when the metallicity is enhanced to 30$\times$ solar."443 Ammonia. and No are expected to be (he most abundant nitrogen-bearing molecules with No becoming, Ammonia and $_2$ are expected to be the most abundant nitrogen-bearing molecules with $_2$ becoming444At;>T.,$\Delta t_i>T$.445" Hence, values of 6P different from zero suggest the existence of correlations between x; and At;."," Hence, values of $\delta P$ different from zero suggest the existence of correlations between $x_i$ and $\Delta446t_i$."447" Evidently, because of statistical fluctuations, 6P is never exactly equal to zero even in catalogs where x; and Af; are uncorrelated."," Evidently, because of statistical fluctuations, $\delta P$ is never exactly equal to zero even in catalogs where $x_i$ and $\Delta t_i$ are uncorrelated."448" In order to explicitly take into account the role of statistical fluctuations, we computed the quantity P*(x;>X|At;«T), defined as P, but in a catalog where flare energies were randomly reshuffled."," In order to explicitly take into account the role of statistical fluctuations, we computed the quantity $P^*(x_i>X \vert \Delta t_i<T)$, defined as $P$, but in a catalog where flare energies were randomly reshuffled."449" The comparison with thereshuffled catalog is the basis of the surrogate data technique (Theiller et al. 1992,,"," The comparison with thereshuffled catalog is the basis of the surrogate data technique (Theiller et al. \cite{The},"450" Schreiber Schmitz 2000)), recently applied to investigate coherent structures in space plasma (Sahraoui 1))"," Schreiber Schmitz \cite{Sch}) ), recently applied to investigate coherent structures in space plasma (Sahraoui \cite{Sah}) )."451 In the reshuffled catalog x; was by construction uncorrelated to Af; and P*(x;>X|At;«T) fluctuated around its average value P(x;>X)., In the reshuffled catalog $x_i$ was by construction uncorrelated to $\Delta t_i$ and $P^*(x_i>X \vert \Delta t_i<T)$ fluctuated around its average value $P(x_i>X)$.452" The amplitude of these fluctuations defined the significance level, c(X,T), which allowed us to distiguish between the presence and absence of correlations."," The amplitude of these fluctuations defined the significance level, $\sigma (X,T)$, which allowed us to distiguish between the presence and absence of correlations."453" The method is schematically presented in Fig.1 for x;=E;, X=E3Eo and T= 1h. P'(E;>E|At;«T) takes different values for each realization of the reshuffled catalog."," The method is schematically presented in Fig.1 for $x_i=E_i$, $X=E=3 E_0$ and $T=1$ h. $P^*(E_i>E \vert \Delta t_i<T)$ takes different values for each realization of the reshuffled catalog."454" We produced 10? independent realizations of the catalog with reshuffled energies and observed that P*(E;>E|At;«T) is Gaussian-distributed with a mean P(E;>E) and a standard deviation o(E,T)."," We produced $10^5$ independent realizations of the catalog with reshuffled energies and observed that $P^*(E_i>E \vert \Delta t_i<T)$ is Gaussian-distributed with a mean $P(E_i>E)$ and a standard deviation $\sigma (E,T)$."455 Similar results were obtained for other values of E and T and for the other definitions of αι., Similar results were obtained for other values of $E$ and $T$ and for the other definitions of $x_i$.456" Therefore, |6P(x;>X|At;« indicates that energies in the real catalog follow a significantly different organization than events in the reshuffled catalog."," Therefore, $\vert\delta P(x_i>X \vert \Delta t_i<T)\vert>\sigma (X,T)$ indicates that energies in the real catalog follow a significantly different organization than events in the reshuffled catalog."457 In Fig.2a we plot 6P(E;>E|At;«T) and in Fig.2b óP(E;.|>E|At;«T) for different values of T., In Fig.2a we plot $\delta P(E_i>E \vert \Delta t_i<T)$ and in Fig.2b $\delta P(E_{i-1}>E \vert \Delta t_i<T)$ for different values of $T$.458 The standard deviation for each data point σ(Ε.T) is represented as the error bar.," The standard deviation for each data point $\sigma (E,T)$ is represented as the error bar."459 In Fig.2a we notice that óP(E;>E|At;«T) takes always positive values beyond error bars., In Fig.2a we notice that $\delta P(E_i>E \vert \Delta t_i<T)$ takes always positive values beyond error bars.460 This implies that in the real catalog the number of couples fulfilling both conditions is greater than in catalogs where energies and intertimes are uncorrelated., This implies that in the real catalog the number of couples fulfilling both conditions is greater than in catalogs where energies and intertimes are uncorrelated.461" More precisely we find that for each given value of E, óP(E;>E|At;«T) decreases by increasing T."," More precisely we find that for each given value of $E$, $\delta P(E_i>E \vert \Delta t_i<T)$ decreases by increasing $T$."462 This implies that the probability to find flare couples with the second flare energy higher than E decreases if one includes events with larger At in the analysis., This implies that the probability to find flare couples with the second flare energy higher than $E$ decreases if one includes events with larger $\Delta t$ in the analysis.463" This result is disagrees with what is expected according to the RV model, which predicts larger flares after longer waiting times."," This result is disagrees with what is expected according to the RV model, which predicts larger flares after longer waiting times."464" On the other hand, obscuration can represent a possible explanation of the above result."," On the other hand, obscuration can represent a possible explanation of the above result."465" Indeed, according to the selection procedure, for flares close in time, the second event is recorded in the catalog only if it produces an increase in the flux of at least 40% of the level of the previous flare."," Indeed, according to the selection procedure, for flares close in time, the second event is recorded in the catalog only if it produces an increase in the flux of at least $40\%$ of the level of the previous flare."466" Moreover, the light curve of large flares may hide smallerflares close in time."," Moreover, the light curve of large flares may hide smallerflares close in time."467 Obscuration effects then reduce the probability to find a small flare after a short waiting time and may introduce spurious correlations between waiting time and the successive flare energy., Obscuration effects then reduce the probability to find a small flare after a short waiting time and may introduce spurious correlations between waiting time and the successive flare energy.468 We then explicitly explored the role of obscuration on óP(E;>E|At;« T)., We then explicitly explored the role of obscuration on $\delta P(E_i>E \vert \Delta t_i<T)$ .469 As a first analysis we considered only, As a first analysis we considered only470of the 2;=5 LBOs. we have also calculated stellar mass estimates based on the stacked LBG photometry.,"of the $z\geq5$ LBGs, we have also calculated stellar mass estimates based on the stacked LBG photometry."471 The obvious advantage of this is that the extra depth provided by stacking produces much more robust photometry in the ./ and A bands. and also provides a detection at 3.6//m from the stacked SWIRE data.," The obvious advantage of this is that the extra depth provided by stacking produces much more robust photometry in the $J$ and $K-$ bands, and also provides a detection at $\mu$ m from the stacked SWIRE data."472 Using a straightforward average stack of the LBG data (see Figs 2 3) our SED fit returns a stellar mass estimate of &AL. ., Using a straightforward average stack of the LBG data (see Figs 2 3) our SED fit returns a stellar mass estimate of $\simeq5\times10^{10}\Msolar$ .473 However. this stack is biased due to the inclusion of the two LBG candidates with the lowest stellar mass estimates (MCD| MCD8). which have the faintest /|A photometry.," However, this stack is biased due to the inclusion of the two LBG candidates with the lowest stellar mass estimates (MCD1 MCD8), which have the faintest $J+K$ photometry."474 This is confirmed by an SED fit to a median stack of the LBG photometry. which returns a stellar mass estimate of =1107M. (as anticipated from the results in Table 2).," This is confirmed by an SED fit to a median stack of the LBG photometry, which returns a stellar mass estimate of $\simeq1\times10^{11}\Msolar$ (as anticipated from the results in Table 2)."475 In conclusion. taking into account systematic differences due to the choice of IMF. the evidence from stacking the LBG data suggests a typical stellar mass of 5.LOMM...," In conclusion, taking into account systematic differences due to the choice of IMF, the evidence from stacking the LBG data suggests a typical stellar mass of $\gtsim\,5\times10^{10}\Msolar$."476 The results presented in the previous section suggest that. despite the considerable uncertainties. the typical stellar mass of the 2=5 LBG candidates is 755.1077AL...," The results presented in the previous section suggest that, despite the considerable uncertainties, the typical stellar mass of the $z\geq5$ LBG candidates is $\gtsim\,5\times10^{10}\Msolar$."477 In this section we investigate whether the existence of such massive galaxies at this early epoch is consistent with ACDM structure formation and current galaxy formation models., In this section we investigate whether the existence of such massive galaxies at this early epoch is consistent with $\Lambda$ CDM structure formation and current galaxy formation models.478 Within this section we adopt the best-fitting stellar mass estimates for each LBG candidate. based on a Salpeter IMF.," Within this section we adopt the best-fitting stellar mass estimates for each LBG candidate, based on a Salpeter IMF."479 However. when necessary. we adjust our estimated number densities to account for differences due to specific choices of IMF.," However, when necessary, we adjust our estimated number densities to account for differences due to specific choices of IMF."480 The first. model independent. quantity of interest to calculate is the surface density of 2°<25 galaxies at +5.," The first, model independent, quantity of interest to calculate is the surface density of $z^{\prime}\leq25$ galaxies at $z\geq5$."481 However. in order to accurately compute the surface density. it is first necessary to account for the inevitable incompleteness introduced during the construction of the original SExtractor (Bertin Arnouts 1996) catalogues due to object blending.," However, in order to accurately compute the surface density, it is first necessary to account for the inevitable incompleteness introduced during the construction of the original SExtractor (Bertin Arnouts 1996) catalogues due to object blending."482 To quantify this effect we performed simulations based on introducing 1000 fake LBG candidates at a time into the τ΄ band images. and attempting to recover them using the same SExtractor configuration adopted for constructing the original σ΄ band catalogues.," To quantify this effect we performed simulations based on introducing 1000 fake LBG candidates at a time into the $z^{\prime}-$ band images, and attempting to recover them using the same SExtractor configuration adopted for constructing the original $z^{\prime}-$ band catalogues."483 This immediately revealed that on average we lose 10% of possible LBG candidates due to object blending on the σ΄ band images., This immediately revealed that on average we lose $\simeq10\%$ of possible LBG candidates due to object blending on the $z^{\prime}-$ band images.484 In addition. we then proceeded to run SExtractor in two-image mode on the corresponding 2 and V band data. using the z/ band as the detection image.," In addition, we then proceeded to run SExtractor in two-image mode on the corresponding $B$ and $V-$ band data, using the $z^{\prime}-$ band as the detection image."485 This revealed that we lose further LBG candidates due to contamination of the matched apertures in the /5 and V bands by nearby companions., This revealed that we lose further LBG candidates due to contamination of the matched apertures in the $B$ and $V-$ bands by nearby companions.486 In around 75'4 of cases this contamination is sufficient to fail our selection criteria requiring <27 detections in both the £3 and V. bands., In around $\simeq5\%$ of cases this contamination is sufficient to fail our selection criteria requiring $\leq2\sigma$ detections in both the $B$ and $V-$ bands.487 Consequently. when calculating the surface density of z<25 LBGs it is necessary to scale-up by a factor of &1.2.," Consequently, when calculating the surface density of $z^{\prime}\leq25$ LBGs it is necessary to scale-up by a factor of $\simeq1.2$."488" Allowing for this correction factor, our estimate for the surface density of Vox25 LBGs at z>5 is 0.005+0.002 per square aremin."," Allowing for this correction factor, our estimate for the surface density of $z^{\prime}\leq25$ LBGs at $z\geq5$ is $0.005\pm0.002$ per square arcmin."489 We note that this figure is compatible with. although obviously much more robust than. a prediction based on the one previously known >5 LBG with 2x5525 in a compilation of deep HST survey fields: 0.004= per square aremin (Bouwens et al.," We note that this figure is compatible with, although obviously much more robust than, a prediction based on the one previously known $z\geq5$ LBG with $z_{850}\leq25$ in a compilation of deep HST survey fields; $0.004\pm0.004$ per square arcmin (Bouwens et al."490 2006)., 2006).491 We also note that our final list of nine <<25 LBGs at z>5 is fully consistent with the previous findings of Shimasaku et al. (, We also note that our final list of nine $z^{\prime}\leq25$ LBGs at $z\geq5$ is fully consistent with the previous findings of Shimasaku et al. (4922005) in the SDF.,2005) in the SDF.493 Within the 767 square arcmin area of the SDF. Shimasaku et al.," Within the 767 square arcmin area of the SDF, Shimasaku et al."494 did not tind any possible 2’<25 LBOGs. but did identify 12 candidates in the magnitude range 25.4<z/26.6.," did not find any possible $z^{\prime}\leq25$ LBGs, but did identify 12 candidates in the magnitude range $25.4<z^{\prime}\leq26.6$."495 However. it is important to remember that the selection technique adopted by Shimasaku et al.," However, it is important to remember that the selection technique adopted by Shimasaku et al."496 was tuned to identify LBGs in the redshift interval 5.6<z«6.2., was tuned to identify LBGs in the redshift interval $5.6<z<6.2$.497 Only one of our final list of nine candidates lies within this redshift range (MCDI., Only one of our final list of nine candidates lies within this redshift range (MCD1).498 Consequently. based on our final sample. we would predict only ~0.4 LBG candidates with ο)τς25 within the SDF area which would satisfy the Shimasakuet al.," Consequently, based on our final sample, we would predict only $\simeq0.4$ LBG candidates with $z^{\prime}\leq25$ within the SDF area which would satisfy the Shimasakuet al."499 selection criteria. entirely consistent with their tinding of none.," selection criteria, entirely consistent with their finding of none."500 The best-fitting stellar mass estimates listed in Table 2 suggest tha five of the LBG candidates have masses of M2:101ML...," The best-fitting stellar mass estimates listed in Table 2 suggest that five of the LBG candidates have masses of M $\gtsim\,10^{11}\Msolar$."501" In order to convert the estimated surface density of M ΙΟΥΝ. LBGs a >25 into anumber density. it is necessary to compute the effective co-moving volume of our survey (defined by the survey area and the redshift intervals within which each of the LBG candidates coule have been detected: 2,7,«2€ Sane)."," In order to convert the estimated surface density of M $\gtsim\,10^{11}\Msolar$ LBGs at $z\geq5$ into a number density, it is necessary to compute the effective co-moving volume of our survey (defined by the survey area and the redshift intervals within which each of the LBG candidates could have been detected; $z_{min}<z<z_{max}$ )."502 Due to our adopted optica selection criteria. our selection function has a sharp low-redshifi cut-off at εν=5.," Due to our adopted optical selection criteria, our selection function has a sharp low-redshift cut-off at $z_{min}=5$."503 To compute ρω we have taken the SED fit to each candidate and redshifted it (keeping age. reddening anc mass constant) until the candidate fades below our magnitude limi (= 25).," To compute $z_{max}$ we have taken the SED fit to each candidate and redshifted it (keeping age, reddening and mass constant) until the candidate fades below our magnitude limit $z^{\prime}=25$ )."504 This calculation reveals that. for the tive LBG candidates with M2:1044AL. . the effective volume of the survey is equivalen," This calculation reveals that, for the five LBG candidates with M $\gtsim\,10^{11}\Msolar$ , the effective volume of the survey is equivalent"505with the fact there have been no reports of the detection of a pulse in the lower energy baud for this pulsar.,with the fact there have been no reports of the detection of a pulse in the lower energy band for this pulsar.506 The 5-ray light curve from Ferm: (Weltevredeetal.2010) shows a broad. peak at ó0.2-," The $\gamma$ -ray light curve from \citep{We10}507 shows a broad peak at $\phi\sim0.2$ -0.5."508 This peak may consist of two components. but it is uot clear iu the current plioton statistics.," This peak may consist of two components, but it is not clear in the current photon statistics."509 The X-ray. pulse profile from is cletected weakly at a mareinal level. aud shows a broad peak at ο 0.6-0.7. which is different. [rom tlie s-ray peak (Roberts.Romani&Johuston 2009).," The X-ray pulse profile from is detected weakly at a marginal level, and shows a broad peak at $\phi\sim0.6$ -0.7, which is different from the $\gamma$ -ray peak \citep{Ro01, B09}."510. Recently. iu the table of Marelli.DeLuca&Caraveo(2011) they list this object as nou-therma| dominated source in X-ray.," Recently, in the table of \citet{MDC11} they list this object as non-thermal dominated source in X-ray."511 The pulsed X-ray. profile is likely to originate from the uou-thermal compouent., The pulsed X-ray profile is likely to originate from the non-thermal component.512 Gj=nee the light curve of this pulsar anc its geometrical parameters are similar to those of PSR JOGSI+1HL1. we adopt the same interpretation.," Since the light curve of this pulsar and its geometrical parameters are similar to those of PSR J0659+1414, we adopt the same interpretation."513 That is. tle ?-ray peak is formecl by outward eluission. whereas the X-ray peak ts formed by inward emission.," That is, the $\gamma$ -ray peak is formed by outward emission, whereas the X-ray peak is formed by inward emission."514 From the intensity map in Fie. G((, From the intensity map in Fig. \ref{fig:3.1}( (515E). an eluission altitude of rj; 1.10-1.11 corresponds to oue broad peak at ó~0.2—0.5 by outward eimmission aud anotler al ᾧ0.6—0.7 by inward emission.,"E), an emission altitude of $r_{ov}\sim1.10$ -1.11 corresponds to one broad peak at $\phi\sim0.2-0.5$ by outward emission and another at $\phi\sim0.6-0.7$ by inward emission."516 Here a sinall shift 96=0.10 toward earlier phase is used., Here a small shift $\delta\phi=0.10$ toward earlier phase is used.517 Since there are similarities in both the 2- and. X-ray. ligit curves aud the geometrical parameters between this pulsar aud PSR J06594-1LLL. we expect a simular double-peak pulse profile iu the optical baud. if it is detected.," Since there are similarities in both the $\gamma$ - and X-ray light curves and the geometrical parameters between this pulsar and PSR J0659+1414, we expect a similar double-peak pulse profile in the optical band, if it is detected."518 Observations iu the 7-ray band liave been obtaited byFerme (Abdoetal.2009d) andACILE (Halpernetal. 2008).., Observations in the $\gamma$ -ray band have been obtained by \citep{Ab20} and \citep{Ha08}. .519 TIe observed light curve slws a sharp double-peak structure., The observed light curve shows a sharp double-peak structure.520 The first peak is ollset [rom the ralio peak by ó~0.16 aud tle two peaks are separated by Ao~0.17., The first peak is offset from the radio peak by $\phi\sim0.16$ and the two peaks are separated by $\Delta\phi\sim0.47$.521 The X-ray light cu've in Abdoetal.(2009d) shows a reatively sharp peak associated with first peak in the s-ray light curve alelt with weak photon staistics., The X-ray light curve in \citet{Ab20} shows a relatively sharp peak associated with first peak in the $\gamma$ -ray light curve albeit with weak photon statistics.522 In this paper. we assume that at least the first peak 1s Dou-therimal in origin.," In this paper, we assume that at least the first peak is non-thermal in origin."523 The possible coutribution of non-thermal X-ray. ejulssions is also cliscussed iu Hesselseal.(2001). and VanEtte1.Romani&Ng(2008)., The possible contribution of non-thermal X-ray emissions is also discussed in \citet{He04} and \citet{VRN08}.524. We expect that this assumption wiIL be tested wy phase-resolved spectra vow future observations., We expect that this assumption will be tested by phase-resolved spectra from future observations.525 As seen i the upper pauel of Fig. 6((, As seen in the upper panel of Fig. \ref{fig:3.1}( (526F). the enission altitude is rj; 0.97-0.98. for which there is a 5-ray clouble-peak profile with separation Ao=0.17 and a relative shift 60=0.06 toward later pase.,"F), the emission altitude is $r_{ov}\sim0.97$ -0.98, for which there is a $\gamma$ -ray double-peak profile with separation $\Delta\phi=0.47$ and a relative shift $\delta \phi=0.06$ toward later phase."527 The peak of non-thermal emission at o= 0.15-0.20 in the S-ray light. curve is found to be fo1ued by outward emission ouly., The peak of non-thermal emission at $\phi=0.15$ -0.20 in the X-ray light curve is found to be formed by outward emission only.528 The relatively weak secoud peak iu the X-ray. band is consistent. with the case of PSRJO205+6119., The relatively weak second peak in the X-ray band is consistent with the case of PSRJ0205+6449.529 Thus. the thiree model parameters for this pulsar audPSR JO205+6119 are very similar. as shown in Table 1.," Thus, the three model parameters for this pulsar andPSR J0205+6449 are very similar, as shown in Table 1."530When |5]2»1. we seek for a WIND solution of Eq. (8)).,"When $\vert \gamma \vert \gg 1$, we seek for a WKB solution of Eq. \ref{EquationMHD}) )."531 Up to second order in e—1/5. a solution for the homogeneous part can be found as: llere. Al and © are two integration constants: Wn) =7+ − ICE + NIE Il(7)) = ((1-," Up to second order in $\epsilon = 1/\gamma$, a solution for the homogeneous part can be found as: Here, $A$ and $\phi$ are two integration constants: ) = + - ) + ) , ) = (1 -."532" Using the method of variation of constants. one obtains the solution for the magnetic potential ancl the vorticity: ((r))= [dt i Τε ? [ism UT) 2o= di2( J) « To """," Using the method of variation of constants, one obtains the solution for the magnetic potential and the vorticity: ) = dt - ) + ^2 ], ) = dt - ) + ^2 ]"533" Using the method of variation of constants. one obtains the solution for the magnetic potential ancl the vorticity: ((r))= [dt i Τε ? [ism UT) 2o= di2( J) « To ""m"," Using the method of variation of constants, one obtains the solution for the magnetic potential and the vorticity: ) = dt - ) + ^2 ], ) = dt - ) + ^2 ]"534SUL. the results of our modeling give considerable creclence to the interpretation that the blueshift observed in WD 209453bs transmission specirunm results [rom high altitude winds. as our models produce Doppler shifts of exactly Chat size.,"Still, the results of our modeling give considerable credence to the interpretation that the blueshift observed in HD 209458b's transmission spectrum results from high altitude winds, as our models produce Doppler shifts of exactly that size."535 An intriguing alternate interpretation of the ?. observed blueshift has been proposed by 2.. who show that a small eccenlricily in the orbit of ILD. 209453b can also produce a net blueshift in the planet's óransmission spectrum from orbital motion alone.," An intriguing alternate interpretation of the \citet{sne10} observed blueshift has been proposed by \citet{mon11}, who show that a small eccentricity in the orbit of HD 209458b can also produce a net blueshift in the planet's transmission spectrum from orbital motion alone."536 The eccentricity of ILD 209453bs. orbit is strongly constrained ancl consistent. wilh zero (??)..," The eccentricity of HD 209458b's orbit is strongly constrained and consistent with zero \citep{win05, dem05}."537 However. using the 3-0 upper limits on the planets eccentricity. ? find that a blueshilt of up to 1 kins 1 could result. from a non-cireular orbit. which is consistent with the ?./— observed. velocity shift within its error bars.," However, using the $\sigma$ upper limits on the planet's eccentricity, \citet{mon11} find that a blueshift of up to 1 km $^{-1}$ could result from a non-circular orbit, which is consistent with the \citet{sne10} observed velocity shift within its error bars."538 Ultimately. better constraints on the orbital eccentricity of ID 209458b will be neecled to resolve whether winds or a non-circular orbit are the cause for the observed blueshift in the planet's spectrum.," Ultimately, better constraints on the orbital eccentricity of HD 209458b will be needed to resolve whether winds or a non-circular orbit are the cause for the observed blueshift in the planet's spectrum."539 In the meantime. a probable very low eccentricity for ILD 209458b does require that winds play the dominant role in producing km ! blueshilts in the planet's Uransmission spectrum.," In the meantime, a probable very low eccentricity for HD 209458b does require that winds play the dominant role in producing km $^{-1}$ blueshifts in the planet's transmission spectrum."540 Even more ambitious observations aimed αἱ spatially resolving exoplanet winds bv observing Doppler shifts for individual spectral lines ancl Doppler shifts as a function of transit phase will likely remain bevond (he reach of current observational facilities., Even more ambitious observations aimed at spatially resolving exoplanet winds by observing Doppler shifts for individual spectral lines and Doppler shifts as a function of transit phase will likely remain beyond the reach of current observational facilities.541 Future instrumentation — high resolution spectrographs with A>10? on next generation 30-m class telescopes may be able to measure some of these effects., Future instrumentation – high resolution spectrographs with $R \gtrsim 10^5$ on next generation 30-m class telescopes – may be able to measure some of these effects.542 A combination of vertical wind shear measurements obtained from the Doppler shifts of individual spectral lines along with ingress and egress measurements of wind speeds on opposing limbs of the planet can produce a 3-D “imap” of winds along an exoplanet's terminator., A combination of vertical wind shear measurements obtained from the Doppler shifts of individual spectral lines along with ingress and egress measurements of wind speeds on opposing limbs of the planet can produce a 3-D “map” of winds along an exoplanet's terminator.543 This is of particular interest for tidallv locked hot Jupiters. since the winds at the terminator are intimately Ged to the dav-to-night heat flow and therefore to the global energy budget of the planet.," This is of particular interest for tidally locked hot Jupiters, since the winds at the terminator are intimately tied to the day-to-night heat flow and therefore to the global energy budget of the planet."544 We caution that the form for magnetic drag used in these 2-D models is simplistic and a more realistic (reatmment could result in a different [low pattern. which would alter the details of the predicted ineress/eeress and vertical shear measurements.," We caution that the form for magnetic drag used in these 3-D models is simplistic and a more realistic treatment could result in a different flow pattern, which would alter the details of the predicted ingress/egress and vertical shear measurements."545 However. wilh these models we have been able to demonstrate the twpes of observations that would constrain the 3-D atmospheric circulation.," However, with these models we have been able to demonstrate the types of observations that would constrain the 3-D atmospheric circulation."546 Regardless of the exact form lor (he magnete drag. what a fairly robust result is that the drag should work to reduce wind speeds. resulting in smaller Doppler shifts in the observed spectra.," Regardless of the exact form for the magnetic drag, what a fairly robust result is that the drag should work to reduce wind speeds, resulting in smaller Doppler shifts in the observed spectra."547 In a particularly interesting proof of concept. ?. observed the transmission spectrum of Venus during its 2004 transit of the Sun and discovered Doppler shilted spectral line proliles.," In a particularly interesting proof of concept, \citet{hed11} observed the transmission spectrum of Venus during its 2004 transit of the Sun and discovered Doppler shifted spectral line profiles."548" While many differences exist between. Venus and hot Jupiters. the very. slow (retrograde) rotalion of Venus means (hat it has verv long days (70.5 Venus vears). compared to hot Jupiters"" permanent davs. resulting in some similarities between their circulation regimes."," While many differences exist between Venus and hot Jupiters, the very slow (retrograde) rotation of Venus means that it has very long days $\sim$ 0.5 Venus years), compared to hot Jupiters' permanent days, resulting in some similarities between their circulation regimes."549many of the model combinations cousidered by Brownieetal. (2003).. the difference in UV. field is too low to account for the corresponding difference iu temperature.,"many of the model combinations considered by \cite{BTS}, the difference in UV field is too low to account for the corresponding difference in temperature."550 Tn particular. the combinations proposed for the nearby diffuse interstellar medi require an additioual heating source iu the warm. component.," In particular, the combinations proposed for the nearby diffuse interstellar medium require an additional heating source in the warm component."551" For UW. observations towards Galactic stars. some of the absorption may arise frou, gas in the vicinity of the star (e.g.Doisséetal. 2005]."," For $_2$ observations towards Galactic stars, some of the absorption may arise from gas in the vicinity of the star \citep[e.g.][]{Boisse05}. ."552. In a conrpalioln paper (Nelinéoetal.2008).. we presented a multi-vaveleusth study conibinius spectroscopic UV. optical. IR aud radio observations of the interstellar matter. along the line of sight to the nearby (170 pc) moderately reddened (E(CD-V) = 0.1 star ΠΙΟ 102065.," In a companion paper \citep{NCI}, we presented a multi-wavelength study combining spectroscopic UV, optical, IR and radio observations of the interstellar matter, along the line of sight to the nearby (170 pc) moderately reddened (E(B-V) = 0.17) star HD 102065."553 Absorption observations provide the column deusities of IT» in he J=0 to 5 states. 13n its three fine structure states. CO in the J=0 to 2 states. CTL aud CTL! .," Absorption observations provide the column densities of ${\rm {H}_{2}}$ in the J=0 to 5 states, in its three fine structure states, ${\rm CO}$ in the J=0 to 2 states, ${\rm CH}$ and ${\rm CH}^{+}$ ."554 They are complemented by. observations ofthe C. COWTy and (10) lines aud dust contiuuua in chussion.," They are complemented by observations of the ${\rm C}^{+}$ , ${\rm CO(2-1)}$ and $(1-0)$ lines and dust continuum in emission."555 Nou-cdetections of the C5 am CN optical absorption lues were transformed iuto wpper lnüts on cohuun densities using oscillator streneths listed by Credeletal.(1991. 1993).," Non-detections of the $_2$ and CN optical absorption lines were transformed into upper limits on column densities using oscillator strengths listed by \cite{Gredel91,Gredel93}."556". The TD 102065 line of sight is well suited for detailed modeling of plysical conditions. cliemustry aud II, excitation. because of the large amount of available data (Table 1))."," The HD 102065 line of sight is well suited for detailed modeling of physical conditions, chemistry and $_2$ excitation, because of the large amount of available data (Table \ref{tab:XmodXobs}) )."557 The molecular fraction is not altered by he preseuce of warn atomic gas along the line of sight., The molecular fraction is not altered by the presence of warm atomic gas along the line of sight.558 Comparison of II T. IT) aud dust extinction indicates that he bulk }) of the coli density is accounted for by a diffuse molecular cloud ideutified ou IRRASimages.," Comparison of H I, $_2$ and dust extinction indicates that the bulk ( ) of the column density is accounted for by a diffuse molecular cloud identified on IRASimages."559 RAS observations place a stroug coustraiut ou the presence of natter close to the star., IRAS observations place a strong constraint on the presence of matter close to the star.560 The abundance of CIT! is hieh., The abundance of $^+$ is high.561 This paper exteuds the work of Cavotal.(2002).. where oulv the FUSE II» spectra of TD 102065 was analyzed. O0 a wider set of observations. usinge an updated version : the Meudon PDR model (LePetitetal.," This paper extends the work of \cite{GBNPHF}, where only the FUSE $_2$ spectrum of HD 102065 was analyzed, to a wider set of observations, using an updated version of the Meudon PDR model \citep{LNHLR}."562.2006).. Πορ) used this sight line to discuss the impact of παν extinction ou the CO abuucdance., \cite{Kopp00} used this sight line to discuss the impact of far-UV extinction on the CO abundance.563 The structure of the paper is as follow: Sect., The structure of the paper is as follow: Sect.564 2 preseuts ιο PDR model used to characterize the diffuse molecular cloud and Sect., \ref{A-simple-PDR} presents the PDR model used to characterize the diffuse molecular cloud and Sect.565 3. describes the main modeling results.," \ref{sec:observationsmodelling}566 describes the main modeling results."567 First. a reference model is defined by fitting observational constraints together.," First, a reference model is defined by fitting observational constraints together."568 Secoud. the model is compared with each of the observations to assess the depeudence of model predictions on the values of the physical parameters.," Second, the model is compared with each of the observations to assess the dependence of model predictions on the values of the physical parameters."569 Iu Sect. 1.," In Sect. \ref{sec:warmh2},"570 we present a detailed attempt to model IT. excitation with a warmer component located close to the star., we present a detailed attempt to model $_2$ excitation with a warmer component located close to the star.571 Sect., Sect.572 5. prescuts our conclusions., \ref{sec:Conclusions} presents our conclusions.573 We use a comprehensive model of an interstellar cloud. that describes the state of the gas aud dust exposed to a radiation field as a function of optical depth," We use a comprehensive model of an interstellar cloud, that describes the state of the gas and dust exposed to a radiation field as a function of optical depth."574 The model is ouc-cimensional aud stationary., The model is one-dimensional and stationary.575 It. prescuts several miproveimnents over that previously used by etal.(2002). and is described in detail iu LePetitetal.(20063., It presents several improvements over that previously used by \cite{GBNPHF} and is described in detail in \cite{LNHLR}.576. It computes simultaneously. iu an iterative wav: Model parameters (see Table 2)) are kept fixed at values consistent with typical diffuse clouds aud the measured characteristics of the ΠΟ 102065 line of sight.," It computes simultaneously, in an iterative way: Model parameters (see Table \ref{model_param}) ) are kept fixed at values consistent with typical diffuse clouds and the measured characteristics of the HD 102065 line of sight."577 The model ignores the eas velocity structure and assmnues that the hree low velocity components discussed by Neluuóéetal.(2008) iae a single homogencous cloud with a visual extinction equal to the observed value to ΠΟ 102065., The model ignores the gas velocity structure and assumes that the three low velocity components discussed by \cite{NCI} make a single homogeneous cloud with a visual extinction equal to the observed value to HD 102065.578 We ook for a best fi model by varving ouly the UV radiatiou Ποια streugth € (isotropic and incident on both sides ofthe cloud) measured in units of Draine’s radiation feld. aud he gas deusitv ay.," We look for a best fit model by varying only the UV radiation field strength G (isotropic and incident on both sides of the cloud) measured in units of Draine's radiation field, and the gas density $n_{\rm H}$ ."579 Density is constaut throughout the cloud., Density is constant throughout the cloud.580 Temperature is computed by solving for thermal valance., Temperature is computed by solving for thermal balance.581 The standard Draine’s field (G=1. Draine. 1978)) is equivalent to about 1.6 iu units of Tabine’s field (IIabiug.1968).," The standard Draine's field $G=1$, \citealp{Draine78}) ) is equivalent to about 1.6 in units of Habing's field \citep{Habing68}."582. Element abundauces are those measured for ¢ Oplitakenfrom Savage&Sembach, Element abundances are those measured for $\zeta$ Ophtakenfrom \cite{SavSem}. .583 , Fig.584(1996).. Fie. 1 shows that ΠΟ102065 measured abundances are close to these referencevalues (Paper D., \ref{abun-figure} shows that HD102065 measured abundances are close to these referencevalues (Paper I).585 The erain distribution is kept fixed., The grain distribution is kept fixed.586 It is a MRN type one with a power lw size, It is a MRN type one with a power law size587We are grateful to Al Glassgold. Barbara Ercolano. and. David. Arclila [or interesting and useful discussions regarding the interpretation of the observations.,"We are grateful to Al Glassgold, Barbara Ercolano, and David Ardila for interesting and useful discussions regarding the interpretation of the observations."588 JN thanks for (heir eenerous hospitality Tom Soifer and (heSpitzer Science Center. where much of the analysis for Chis paper was carried out.," JN thanks for their generous hospitality Tom Soifer and the Science Center, where much of the analysis for this paper was carried out."589 This work is based on observations made with theTelescope. which is operated by the Jet. Propulsion Laboratory. California Institute ol Technology under a contract with NASA.," This work is based on observations made with the, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA."590 Support for this work was provided by NASA through an award issued by JPL/Caltech., Support for this work was provided by NASA through an award issued by JPL/Caltech.591 Basic research in infrared astronomy at (he Naval Research Laboratory is supported by 6.1 base funding., Basic research in infrared astronomy at the Naval Research Laboratory is supported by 6.1 base funding.592The main characteristics of the model are given in Table 2.,The main characteristics of the model are given in Table \ref{SCBparam}.593 Phe stellar viclels are κο as described: in section 77 for massive ancl intermediate mass stars., The stellar yields are used as described in section \ref{parameters} for massive and intermediate mass stars.594 The stellar lifetimes come from (Polsetal.1997).. ancl are dependent on metallicity.," The stellar lifetimes come from \shortcite{97POL_EA}, and are dependent on metallicity."595 “Phe onset of type la supernovac happens when the metallicity reaches Fe/11]—-1.0., The onset of type Ia supernovae happens when the metallicity reaches [Fe/H]=-1.0.596 The vields for the dillerent species produced by type la SN are from Nomoto et al. C, The yields for the different species produced by type Ia SN are from Nomoto et al. (5971984) (iron z 0.6 AL. fevent) aid the rate of SNla is assumed to be proportional to the number of SNL. with SNIL/SNLIa-8.5.,"1984) (iron $\approx$ 0.6 $_{\odot}$ /event) and the rate of SNIa is assumed to be proportional to the number of SNII, with SNII/SNIa=8.5."598 Note that we don't use a Schmidt [aw tvpe SER., Note that we don't use a Schmidt law type SFR.599 While there are some evidences that the SER. of massive stars may be proportional to some power of the gas density. there are few evidences that this can be extrapolated to low and intermediate mass stars.," While there are some evidences that the SFR of massive stars may be proportional to some power of the gas density, there are few evidences that this can be extrapolated to low and intermediate mass stars."600" Since determinations to date are compatible with a constant SER. history for the disc of the Milky: Way. we use this simpleprescription""."," Since determinations to date are compatible with a constant SFR history for the disc of the Milky Way, we use this simple."601 The delavecl ejection of important quantities of gas from long-lived stars could enhanced. the dilution of metals in the interstellar medium. and permit the production of more stars at intermediate (ος solar) abundance.," The delayed ejection of important quantities of gas from long-lived stars could enhanced the dilution of metals in the interstellar medium, and permit the production of more stars at intermediate (e.g solar) abundance."602 In our tests however. this elfect proved to be very minor. for the reason that the release of gas is spread. on long time-scale by the important variation of stellar lifetime with abundance and mass: therefore. there is no sudden release of gas. ancl the dilution of metals in the interstellar medium. is smoothee over long time scales.," In our tests however, this effect proved to be very minor, for the reason that the release of gas is spread on long time-scale by the important variation of stellar lifetime with abundance and mass; therefore, there is no sudden release of gas, and the dilution of metals in the interstellar medium is smoothed over long time scales."603 Phe only marked effect of the SCB model with no LRA is the more rapid. metallicity increase at the onset of SNla at Fe/11]—-1., The only marked effect of the SCB model with no IRA is the more rapid metallicity increase at the onset of SNIa at [Fe/H]=-1.604 This translates in the metallicity distribution into a shift of the distribution to higher metallicities. as can be seen at. Fe/I]e-1 (Fig.13aa).," This translates in the metallicity distribution into a shift of the distribution to higher metallicities, as can be seen at $\approx$ -1 \ref{closeboxnoira}a a)."605 The SCB with no LRA doesn't seems to generate more solar metallicity stars., The SCB with no IRA doesn't seems to generate more solar metallicity stars.606 Note that this somewhat contracicts the findings of Scully et al. (, Note that this somewhat contradicts the findings of Scully et al. (6071997).,1997).608 However. they use very different prescriptions (stellar lifetimes independent of metallicity. IATL index-1.7 over the whole mass range. which they take to be θες m/M. < 100).," However, they use very different prescriptions (stellar lifetimes independent of metallicity, IMF index=1.7 over the whole mass range, which they take to be $\le$ $_{\odot}$ $\leq$ 100)."609 Vhe total curation of the model is 1H Gyr., The total duration of the model is 14 Gyr.610 In the moclel. we identify three phases corresponding (not necessarily in a univocal correspondance) more or less to the ealactic stellar populations.," In the model, we identify three phases corresponding (not necessarily in a univocal correspondance) more or less to the galactic stellar populations."611 The first phase is defined by Fe/H]«-1.0 dex., The first phase is defined by $<$ -1.0 dex.612 We do not try to ascribe a particular stellar population (halo or thick disc) to this phase. and the scale. height correction attributed to this metallicity range is arbitrary.," We do not try to ascribe a particular stellar population (halo or thick disc) to this phase, and the scale height correction attributed to this metallicity range is arbitrary."613 In other words. we don't include this part of the distribution in our discussion. of the SCB model.," In other words, we don't include this part of the distribution in our discussion of the SCB model."614 This. conservative position is justified to some extent by (1) the fact that the characteristics of the metal weak thick clise (or [attened halo 7), This conservative position is justified to some extent by (1) the fact that the characteristics of the metal weak thick disc (or flattened halo ?)615 are still essentially unknown (2) the limited volume of our, are still essentially unknown (2) the limited volume of our616The interstellar medium (ISM) is à multi-component plasma that consists mostly of hydrogen. helium. heavy tons. electrons and charged dust grains.,"The interstellar medium (ISM) is a multi-component plasma that consists mostly of hydrogen, helium, heavy ions, electrons and charged dust grains."617 The interaction. between these components. and their coupling with the Galactic magnetic field. determine the dynamical properties of the ISM. control its evolution and the nature of the star formation process.," The interaction between these components, and their coupling with the Galactic magnetic field, determine the dynamical properties of the ISM, control its evolution and the nature of the star formation process."618 In particular. the momentum exchange in collisions between neutral and charged particles is responsible for transfering the effects of electric and magnetic forces to the neutral component. allowing the magnetic field to drift out of weakly-ionized molecular clouds (Mestel Spitzer 1956). damping the propagation of Alfvenn waves (Zweibel Josafatsson 1983). and heating the gas by the frictional dissipatior of turbulent energy (Scalo 1977).," In particular, the momentum exchange in collisions between neutral and charged particles is responsible for transfering the effects of electric and magnetic forces to the neutral component, allowing the magnetic field to drift out of weakly-ionized molecular clouds (Mestel Spitzer 1956), damping the propagation of Alfvènn waves (Zweibel Josafatsson 1983), and heating the gas by the frictional dissipation of turbulent energy (Scalo 1977)."619 In a companion paper (Pinto. Galli Bacciotti. 2007. hereafter Paper I). we have derived the equations governing the dynamies of a three-fluid system. reducing the set of equations to a momentum equation for the mean fluid and an evolution equation for the magnetic field. plus two relations for the drift velocities in terms of the mean fluid velocity and the magnetic field.," In a companion paper (Pinto, Galli Bacciotti 2007, hereafter Paper I), we have derived the equations governing the dynamics of a three-fluid system, reducing the set of equations to a momentum equation for the mean fluid and an evolution equation for the magnetic field, plus two relations for the drift velocities in terms of the mean fluid velocity and the magnetic field."620 In this paper. we report on a detailed analysis of collisional rate coefficients involving the most abundant neutral and charged species in the ISM.," In this paper, we report on a detailed analysis of collisional rate coefficients involving the most abundant neutral and charged species in the ISM."621 The paper is organized as follows: in Sect. 2..," The paper is organized as follows: in Sect. \ref{sec_friction},"622 we give the general expression for the friction force and the momentum transfer rate coefficient for elastic collisions. and we obtain an analytical solution for a cross section varying as a power of the relative velocity: in Sect. 3..," we give the general expression for the friction force and the momentum transfer rate coefficient for elastic collisions, and we obtain an analytical solution for a cross section varying as a power of the relative velocity; in Sect. \ref{sec_h2coll},"623 we consider collisions with H» of HCO Hy. H™. and electrons. using available theoretical and/or experimental. determination of the collision. cross section; similarly. in Sect. 4.. ," we consider collisions with $_2$ of $^+$, $_3^+$, $^+$, and electrons, using available theoretical and/or experimental determination of the collision cross section; similarly, in Sect. \ref{sec_hcoll}, ,"624we consider collisions with H of C . and electrons: in Sect. 5..," we consider collisions with H of $^+$, $^+$, and electrons; in Sect. \ref{sec_hecoll},"625 we consider collisions of H and electrons with He: in Sects., we consider collisions of $^+$ and electrons with He; in Sects.626 6 and 7 we consider collisions between charged dust grains and neutral particles. and between charged particles. respectively: in Sect. 8..," \ref{sec_graincoll} and \ref{sec_chargecoll} we consider collisions between charged dust grains and neutral particles, and between charged particles, respectively; in Sect. \ref{sec_anal},"627 we give analytical approximations for our numerical results: finally. 1n Sect. 9..," we give analytical approximations for our numerical results; finally, in Sect. \ref{sec_concl},"628 we summarize our conclusions., we summarize our conclusions.629" The general expression of the momentum acquired per unit time and unit volume (“friction force"") by a particle of species 5 with mass 7/5, and initial velocity v, (test particle”) through collisionswith particles of species s with mass zn, and initial velocity vy (“field particles”) was given by Boltzmann (1896). where f(v,) and f(vy) are the velocity distribution functions of the two species. w, is the velocity of the test particle after the collision. dco/d€ is the differential scattering cross section. and v, is the relative velocity of the particles (before the collision). vel."," The general expression of the momentum acquired per unit time and unit volume (“friction force”) by a particle of species $s$ with mass $m_s$ and initial velocity ${\bf v}_s$ (“test particle”) through collisionswith particles of species $s^\prime$ with mass $m_{s^\prime}$ and initial velocity ${\bf v}_{s^\prime}$ (“field particles”) was given by Boltzmann (1896), where $f({\bf v}_s)$ and $f({\bf v}_{s^\prime})$ are the velocity distribution functions of the two species, ${\bf w}_s$ is the velocity of the test particle after the collision, $d\sigma/d\Omega$ is the differential scattering cross section, and $v_{ss^\prime}$ is the relative velocity of the particles (before the collision), |."630 For elastic collisions. the last term in Eq. (1)).," For elastic collisions, the last term in Eq. \ref{fcoll1}) ),"631 representing the momentum change of the test particle after the collision. can be written as Vo) Va. where Jo=manens+ty) is the reduced mass of the system. is thetransfer cross section. and O Is in the center-of-mass system.," representing the momentum change of the test particle after the collision, can be written as _s) _s), where $\mu_{ss^\prime}=m_s m_{s^\prime}/(m_s+m_{s^\prime})$ is the reduced mass of the system, is the cross section, and $\Theta$ is in the center-of-mass system."632 Eq. (1)), Eq. \ref{fcoll1}) )633" reduces then to the expression νοMew, v4)", reduces then to the expression _s) _s).634 The six integrations of eq. (5)), The six integrations of eq. \ref{fcoll2}) )635 require some care because of the presence of the relative velocity in the integrand., require some care because of the presence of the relative velocity in the integrand.636" When the velocity distribution ofboth species is maxwellian with temperatures 7, and 7. the integrations can be carried out explicitly. as shown first by Langevin (1905). and the result is ας)."," When the velocity distribution ofboth species is maxwellian with temperatures $T_s$ and $T_{s^\prime}$ , the integrations can be carried out explicitly, as shown first by Langevin (1905), and the result is _s),"637atmosphere (Trj 1600. log(g)=3) if the influence of the turbulent eas motion on the dust erains is taken into account.,"atmosphere $_{\rm eff}$ =1600, $\log(g)=3$ ) if the influence of the turbulent gas motion on the dust grains is taken into account."638 We interpret our results such that collisional ionization of the dust erains on its own does not provide. in the first place. an ionization level which is sufficientlv high for magnetic coupling.," We interpret our results such that collisional ionization of the dust grains on its own does not provide, in the first place, an ionization level which is sufficiently high for magnetic coupling."639 A collisionally charged dust phase could. however. (rigger secondary non linear chareing processes such as electron avalanches Chat lead to lightning: In the case that the charging process is rather a charge exchange between colliding dust grains. the larger grains will be charged positively ancl the smaller grains would carry awav (he negative charges.," A collisionally charged dust phase could, however, trigger secondary non linear charging processes such as electron avalanches that lead to lightning: In the case that the charging process is rather a charge exchange between colliding dust grains, the larger grains will be charged positively and the smaller grains would carry away the negative charges."640 This may lead to large scale charge separation due to differential sedimentation and also dillerential response to the turbulent gas drag of grains of dilferent nDiass/size., This may lead to large scale charge separation due to differential sedimentation and also differential response to the turbulent gas drag of grains of different mass/size.641 We suggest that. although thermal gas ionisation may decline in objects across the fully-convective boundary. dust ionisation may take over in the lowest mass objects.," We suggest that although thermal gas ionisation may decline in objects across the fully-convective boundary, dust ionisation may take over in the lowest mass objects."642 The onset of atmospheric dust formation. the cloud depth and its particle characteristics may therefore correlate with the anomalous X-ray and radio emission in atmospheres that are cool. but contain highly charged clouds.," The onset of atmospheric dust formation, the cloud depth and its particle characteristics may therefore correlate with the anomalous X-ray and radio emission in atmospheres that are cool, but contain highly charged clouds."643 In the context of this paper. it may be surprising that only a Iraction of Brown Dwarls seem to show activity in form of X-ray or radio emission and not nearly as much is known for exirasolar planets.," In the context of this paper, it may be surprising that only a fraction of Brown Dwarfs seem to show activity in form of X-ray or radio emission and not nearly as much is known for extrasolar planets."644 The observed intermittency of the N-ravs might be interpreted as sien for an intermittent dust cloud distribution as suggested bv us in earlier works (IHelling et al., The observed intermittency of the X-rays might be interpreted as sign for an intermittent dust cloud distribution as suggested by us in earlier works (Helling et al.645 2004)., 2004).646 Depending on the objects parameter like mass. elective temperature. metallicity. age. (he atmosphere does change and so do the conditions for electrification which we are only beginning to explore.," Depending on the objects parameter like mass, effective temperature, metallicity, age, the atmosphere does change and so do the conditions for electrification which we are only beginning to explore."647 Our first study does suggest that cloud electrification is more pronounced in giant planets compared to the more compact brown clwarl atinospheres., Our first study does suggest that cloud electrification is more pronounced in giant planets compared to the more compact brown dwarf atmospheres.648 llowever. bot brown cwarls and giant. planets are prone to lightning events inside (heir dust clouds.," However, both brown dwarfs and giant planets are prone to lightning events inside their dust clouds."649 Systematic investigation of this are part of our future work., Systematic investigation of this are part of our future work.650Recent observations of Type LA Superuovae (SNe) at redshifts up to i~(0.8 (Riessetal.1998: Perhuutter ct al.,Recent observations of Type IA Supernovae (SNe) at redshifts up to $z\sim 0.8$ \cite{ri98}; Perlmutter et al.651 1999. hereafter P99)) have mace possible classical cosimological tests that require standard caudles. such as the maguitude-redshift relation.," 1999, hereafter \cite{per99}) ) have made possible classical cosmological tests that require standard candles, such as the magnitude-redshift relation."652 The most dramatic result is that these SNeappear dimuuer (bv ~0.2 magnitudes) at lugh redshift than would be predicted in a nou-acceleratiug universe. sueeesting at face value that we live in au accelerating universe.," The most dramatic result is that these SNeappear dimmer (by $\sim 0.2$ magnitudes) at high redshift than would be predicted in a non-accelerating universe, suggesting at face value that we live in an accelerating universe."653 However. other explanations are xossible. includiug the one we cousider here. namely that istant SNe are dimuuer due to extinction by intergalactic ust.," However, other explanations are possible, including the one we consider here, namely that distant SNe are dimmer due to extinction by intergalactic dust."654 As distant standard candles. SNe are sensitive probes of extinction in the intergalactic medimu (IGM).," As distant standard candles, SNe are sensitive probes of extinction in the intergalactic medium (IGM)."655 The istribution of iuatter im the ICAL nav now be modelled accurately in the context of modern cosmology sine lvdvodvuamic simulations (sce ce. Cenetal. 1991: Wernguistetal. 1996: Davéetal. 1999)).," The distribution of matter in the IGM may now be modelled accurately in the context of modern cosmology using hydrodynamic simulations (see e.g., \cite{cen94}; \cite{her96}; \cite{dav99}) )."656 The resulting IGAL is not smooth. but rather traces large-scale structure.," The resulting IGM is not smooth, but rather traces large-scale structure."657 Tf such structure contains not only dark matter and gas but also dust. this would result iu significant wariatious in| Όλο brightuesses due to intervening extinction.," If such structure contains not only dark matter and gas but also dust, this would result in significant variations in SNe brightnesses due to intervening extinction."658 Observationallv. the distribution of Type Ia SNe maguitudes has a very small dispersion (P99)).," Observationally, the distribution of Type Ia SNe magnitudes has a very small dispersion \cite{per99}) )."659 Thus by comparing simulations to the distribution of observed brightucsses. we cau set limits on the amount of dust extinction and possibly constrain its spatial distribution with respect to iutergalactic gas.," Thus by comparing simulations to the distribution of observed brightnesses, we can set limits on the amount of dust extinction and possibly constrain its spatial distribution with respect to intergalactic gas."660 In thisLetter we present a techuique for doing this. aud apply it to the SNe observations of P99..," In this we present a technique for doing this, and apply it to the SNe observations of \cite{per99}."661 lutergalactic ervey dust has been examined ina series of papers by Aguirre (1999a.b: hereafter A99)) and Aguirre&Tainan (1999).. who develop a scenario iu which siuall erains are pretercutially destroved during ejection from ealaxies. polluting the [GAL with large dust eraius that are οπουνο exey in the baudpasses of the SNe data.," Intergalactic grey dust has been examined in a series of papers by Aguirre (1999a,b; hereafter \cite{agu99}) ) and \cite{agu99c}, who develop a scenario in which small grains are preferentially destroyed during ejection from galaxies, polluting the IGM with large dust grains that are effectively grey in the bandpasses of the SNe data."662 This erevuess is necessary du order not to violate tieht limits on reddening from SNe data. which Πρίν that ealactic-type dust would provide ueeligible absorption (P99)).," This greyness is necessary in order not to violate tight limits on reddening from SNe data, which imply that galactic-type dust would provide negligible absorption \cite{per99}) )."663 Furthenuore. a significant fraction of the dust iust reside in the IGM.," Furthermore, a significant fraction of the dust must reside in the IGM."664 If the erey cust causing extinction were present only in the ISAL of the supernova host galaxy. this would introduce too large a dispersion in observed SNe magnitudes (Riessotal. 1998)).," If the grey dust causing extinction were present only in the ISM of the supernova host galaxy, this would introduce too large a dispersion in observed SNe magnitudes \cite{ri98}) )."665 In this study. we asstune that erev dust blends sioothhy from galaxies iuto the surroundiug ICAL.," In this study, we assume that grey dust blends smoothly from galaxies into the surrounding IGM."666 Our analysis is insensitive to the intrinsic properties of the dust. such as its opacity and erain size. since we use the sinulatious to directly translate the observed SNe magnitude distribution iuto a dust extinction im maenitucdes.," Our analysis is insensitive to the intrinsic properties of the dust, such as its opacity and grain size, since we use the simulations to directly translate the observed SNe magnitude distribution into a dust extinction in magnitudes."667" It is. however. scusitive to the wav in which dust traces the distribution of eas in the ΤΟΝΙ, and we will consider several simple but plausible variatious of this relation."," It is, however, sensitive to the way in which dust traces the distribution of gas in the IGM, and we will consider several simple but plausible variations of this relation."668 lu& ?? we describe our simulations of the IGML aud of dust extinction.," In \ref{sec: sims} we describe our simulations of the IGM, and of dust extinction."669 Du ?? we describe our analysis method and results. includiue coustraiuts ou erev dust afforded by current SN observations.," In \ref{sec: anal} we describe our analysis method and results, including constraints on grey dust afforded by current SN observations."670 In ο)sve discuss svsteinaticuncertainties. aud the implications of our results.," In \ref{sec: disc} we discuss systematicuncertainties, and the implications of our results."671" We ocuplov ai lydrodvuamic simulation of a A- cold dark matter model. with ο=0. Q4= 0.6. Q,=0.025 7. Hy=OSlars+Mpe and ay= 0.5."," We employ a hydrodynamic simulation of a $\Lambda$ -dominated cold dark matter model, with $\Omega_m=0.4$, $\Omega_\Lambda=0.6$ $\Omega_b=0.02h^{-2}$ , $H_0=65 \kmsmpc$ , and $\sigma_8=0.8$ ."672 ↻↿∐⋅↴∖↴↕∐⋯↕⋜↧⊓∪∐↖↽∪↕∏⋯↸∖↕↴∖↴⋅↱⊐∩∕∣↓⋀∖↕⋉⊳↖↖↽↕↑∐ 10h.| kpe spatial resolution. having LLL? dark matter and LLP gas particles.and was evolved from :=19)>0 using Parallel TreeSPITI (Dave.Dubiuski&IHeruquist 1997)).," Our simulation volume is $50 h^{-1}$ Mpc with $10 h^{-1}$ kpc spatial resolution, having $144^3$ dark matter and $144^3$ gas particles,and was evolved from $z=49\rightarrow 0$ using Parallel TreeSPH \cite{ddh97}) )."673"formed in junctions of multiple intergranular lanes, are places in which the concentration of these magnetic features are expected to be favoured and therefore play an important role in the dynamics of the quiet Sun (Kitiashvilietal.2010).","formed in junctions of multiple intergranular lanes, are places in which the concentration of these magnetic features are expected to be favoured and therefore play an important role in the dynamics of the quiet Sun \citep{kitiashvili2010}."674. We search for the location of small-scale swirls by tracking plasma motions in the same region where Bonetetal.(2008) detected small whirpools by following the trajectories of BPs being swallowed by them., We search for the location of small-scale swirls by tracking plasma motions in the same region where \cite{bonet2008} detected small whirpools by following the trajectories of BPs being swallowed by them.675 From 20-min average flow maps we identify events displaying a converging pattern of horizontal velocity vectors towards a central point that correspond to the strong sinks as initially predicted by the numerical simulations., From 20-min average flow maps we identify events displaying a converging pattern of horizontal velocity vectors towards a central point that correspond to the strong sinks as initially predicted by the numerical simulations.676 These events are in all cases detected in the vertices of multiple granules along very intense intergranular dark lanes., These events are in all cases detected in the vertices of multiple granules along very intense intergranular dark lanes.677 This is clearly seen on the average image for the duration of the time series where the low intensity junctions are enhanced (in Fig. 3))., This is clearly seen on the average image for the duration of the time series where the low intensity junctions are enhanced (in Fig. \ref{corchos}) ).678 Data from two time series and common examination allow the comparison of results., Data from two time series and common examination allow the comparison of results.679" A total of 70 1η 51 and 77 in s2 vortices have been detected, resulting in a density of 2.8 x 10? vortices Mm”? and 3.1 x 107? vortices Mm”? for s1 and s2 respectively."," A total of 70 in s1 and 77 in s2 vortices have been detected, resulting in a density of 2.8 $\times$ $^{-2}$ vortices $^{-2}$ and 3.1 $\times$ $^{-2}$ vortices $^{-2}$ for s1 and s2 respectively."680 Averaging over the 20-min window we obtained space-time-density values of 1.4 x 10? vortices Mm? min! and 1.6 x 107? vortices Mm? min'! for s1 and s2 respectively., Averaging over the 20-min window we obtained space-time-density values of 1.4 $\times$ $^{-3}$ vortices $^{-2}$ $^{-1}$ and 1.6 $\times$ $^{-3}$ vortices $^{-2}$ $^{-1}$ for s1 and s2 respectively.681" Values are comparable to the number obtained by Bonetetal.(2008) of 1.8 x 10? vortices Mm? min“! and lower than the density of 3.1 x 10? vortices Mm""? min""! found by (2010)..", Values are comparable to the number obtained by \citet{bonet2008} of 1.8 $\times$ $^{-3}$ vortices $^{-2}$ $^{-1}$ and lower than the space-time-density of 3.1 $\times$ $^{-3}$ vortices $^{-2}$ $^{-1}$ found by \cite{bonet2010}.682 It is necessary to bear in mind that the results from our flow fields are smoothed by the 20-min averages and the size of the employed tracking window., It is necessary to bear in mind that the results from our flow fields are smoothed by the 20-min averages and the size of the employed tracking window.683 Many short-living vortical motions are likely diluted by the LCT temporal average., Many short-living vortical motions are likely diluted by the LCT temporal average.684 BPs are distributed all over the FOV in this region as found by SánchezAlmeidaetal.(2010) from which some describe spiral trajectories whilst being engulfed by downdrafts that Bonetetal.(2008) described as convectively driven vortex flows., BP«s are distributed all over the FOV in this region as found by \cite{sanchez2010} from which some describe spiral trajectories whilst being engulfed by downdrafts that \cite{bonet2008} described as convectively driven vortex flows.685 We have compared the location of our detected vortices with the events discovered by Bonetetal.(2008) in the same solar region., We have compared the location of our detected vortices with the events discovered by \cite{bonet2008} in the same solar region.686" As a result, we have found of coincidences but the remaining corresponds to cases in the near vicinity of our vortices (i.e. less than 2"" from the edge of the circular area enclosing the events in Fig."," As a result, we have found of coincidences but the remaining corresponds to cases in the near vicinity of our vortices (i.e. less than $\arcsec$ from the edge of the circular area enclosing the events in Fig."687 2 and Fig. 3.., \ref{verticalvel} and Fig. \ref{corchos}.688" We find for the series s1 and s2 counterclockwise/clockwise sense of rotation frequency value ratios of 53/47 and 52/48%.,, respectively, which in view of the number of studied events shows no significant difference with equally probable sense of rotation, in agreement with Bonetetal.(2008)., though Bonetetal.(2010) find a significant preference in favor of counterclockwise sense of rotation."," We find for the series s1 and s2 counterclockwise/clockwise sense of rotation frequency value ratios of 53/47 and 52/48, respectively, which in view of the number of studied events shows no significant difference with equally probable sense of rotation, in agreement with \cite{bonet2008}, though \cite{bonet2010} find a significant preference in favor of counterclockwise sense of rotation."689" The latter explain the different results assuming that the vortex rotation is influenced by the latitudinal solar differential rotation, as the observations in Bonetetal.(2010) were done at mid solar latitude, while the observations in Bonetetal.(2008) and, of course, in this paper were done at the equator, where the influence of differential rotation is negligible."," The latter explain the different results assuming that the vortex rotation is influenced by the latitudinal solar differential rotation, as the observations in \cite{bonet2010} were done at mid solar latitude, while the observations in \cite{bonet2008} and, of course, in this paper were done at the equator, where the influence of differential rotation is negligible."690 In terms of dimensions our detected vortices coincide with the values in Bonetetal.(2010) of less than 500 km of radius., In terms of dimensions our detected vortices coincide with the values in \cite{bonet2010} of less than 500 km of radius.691 The majority of our detected vortices exhibit a radius of 241 + 25 km., The majority of our detected vortices exhibit a radius of 241 $\pm$ 25 km.692" Some vortices, however, have a radius over 400 km that might not"," Some vortices, however, have a radius over 400 km that might not"693This is a quadratic equation for ez which has lowest order solution The kinetic growth rate varies as x(1-+σος). CLaaning2001) and averaging e? over this growth rate gives (02)=2.(L1/2)«2μη.,"This is a quadratic equation for $v_{e\perp}^2$ which has lowest order solution The kinetic growth rate varies as $\gamma\propto\left(6941+{\omega _{pe}^2\over\omega _{pp}^2}\cos ^2\theta\right)^{-1}$, \citep{laming01a}695 and averaging $v_{e\perp}^2$ over this growth rate gives $\left<v_{e\perp}^2\right>=v_{i\perp}^2\left(1-1/\sqrt{2}\right)696\omega _{pe}^2L^2/\omega _{pp}^2r_g^2$."697" Thus in principle electron and ion: temperatures can equilibrate∙∙ (assuming− that. (1—1/v2)⋅Lfr,⋅↽≻⋅↽≻>1). ⋉− in the sense that the waves do not saturate before this is achieved."," Thus in principle electron and ion temperatures can equilibrate (assuming that $\left(1-1/\sqrt{2}\right)L^2/r_g^2 > 1$ ), in the sense that the waves do not saturate before this is achieved."698 We use an adaptation of the BLASPHEMER (BLASt Propagation in Highly EMitting code (Laming2001b:Laming&Grun2002.2003:Lwang2003).. which follows the time dependent ionization balance and (temperatures of a Lagrangian plasma parcel as it expands in the solar wind.," We use an adaptation of the BLASPHEMER (BLASt Propagation in Highly EMitting code \citep{laming01b,laming02,laming03b,laming03c}, which follows the time dependent ionization balance and temperatures of a Lagrangian plasma parcel as it expands in the solar wind."699 The density η of ions of element / with charge q is given by where Chong:Cong:ος ave Lhe rates Lor electron impact ionization. radiative recombination and clielectvonic recombination respectively. out of the charge stale q.," The density $n_{iq}$ of ions of element $i$ with charge $q$ is given by where $C_{ion,q}, C_{rr,q}, C_{dr,q}$ are the rates for electron impact ionization, radiative recombination and dielectronic recombination respectively, out of the charge state $q$."700 These rates are the same as (hose used in the recent ionization balance calculations of Mazzottaοἱal.(1998). using subroutines kindly supplied by Dr P. Alazzotta (private communication 2000).," These rates are the same as those used in the recent ionization balance calculations of \citet{mazzotta98}, using subroutines kindly supplied by Dr P. Mazzotta (private communication 2000)."701" The electron density ο, is determined from the condition that the plasma be electrically neutral.", The electron density $n_e$ is determined from the condition that the plasma be electrically neutral.702" The ion and electron temperatures. Lj, and T; are coupled by Coulomb collisions by and"," The ion and electron temperatures, $T_{iq}$ and $T_e$ are coupled by Coulomb collisions by and"703and. also. of the cloud relative to the intercloud. medium.,"and, also, of the cloud relative to the intercloud medium."704 However. NTO09 cid. include in their data analysis any variation of the magnetic field. [rom each. core to its envelope for two reasons: (1) There is no reliable way to deduce from the CLETO9 data the direction of the magnetic Ποιά either in the. cores or in the envelopes. a fact that would introduce significant uncertainties in the analysis. (," However, MT09 did include in their data analysis any variation of the magnetic field from each core to its envelope for two reasons: (1) There is no reliable way to deduce from the CHT09 data the direction of the magnetic field either in the cores or in the envelopes, a fact that would introduce significant uncertainties in the analysis. ("7052) MTOO wanted to demonstrate that. when CIITOO's overly restrictive assumption of spatially constant. value of 1e field. in each envelope is relaxed. as suggested: by the lata themselves. and proper upper limits are reported. for 16 nondetections. the strong conclusion of CIETO9 on the Pvariation of the mass-to-Iux ratio from cores to envelopes Is ga10wn to be overinllated and completely. unjustified.,"2) MT09 wanted to demonstrate that, when CHT09's overly restrictive assumption of spatially constant value of the field in each envelope is relaxed, as suggested by the data themselves, and proper upper limits are reported for the nondetections, the strong conclusion of CHT09 on the variation of the mass-to-flux ratio from cores to envelopes is shown to be overinflated and completely unjustified."706" More recentlv. Cruteher. Hakobian ""Troland. (2010: —iereinafter. CLETIO) responded to the N'TO9 paper by (1) ""laiming that the ALPO9 analvsis is not. self-consistent. in ju a cosine [factor is missing. which would account for the different directions of the field in a core and its envelope: and by (2) presenting additional arguments for which the original CILTOO data analysis is proper."," More recently, Crutcher, Hakobian Troland (2010; hereinafter CHT10) responded to the MT09 paper by (1) claiming that the MT09 analysis is not self-consistent, in that a cosine factor is missing, which would account for the different directions of the field in a core and its envelope; and by (2) presenting additional arguments for which the original CHT09 data analysis is proper."707 As already explained above. the claim about a missing cosine factor is an inaccurate representation of the MEO9 analysis. which did not assume a different direction of the core and envelope magnetic fieldvectors.," As already explained above, the claim about a missing cosine factor is an inaccurate representation of the MT09 analysis, which did not assume a different direction of the core and envelope magnetic field."708" Ne quote from MTO9. end of .22: ""ln our analysis. we relaxed only of the CLT assumptions (that of lack of spatial variation of Boy. which is not consistent with the data)."," We quote from MT09, end of 2: “In our analysis, we relaxed only of the CHT assumptions (that of lack of spatial variation of $B_{\rm env}$, which is not consistent with the data)."709 We have retained the implicit assumption of similar orientations of the DB and Boe (vechlors). because the data do not suggest. any particular relative orientation of the two vectors.," We have retained the implicit assumption of similar orientations of the ${\mathbf B}_{\rm env}$ and ${\mathbf B}_{\rm core}$ ), because the data do not suggest any particular relative orientation of the two vectors."710 A more general analysis that would also relax this assumption would increase the uncertainties on Z? (although not on D.) and would further part from the CLIEE conclusions.”, A more general analysis that would also relax this assumption would increase the uncertainties on $R$ (although not on $B_{\rm env}$ ) and would further part from the CHT conclusions.”711 The main point of the MTO9 paper is that the CIUTOO data analysis is seriously [awed. even if one ignores the theoretically expected possibility that the core and envelope magnetic fields will have different directions., The main point of the MT09 paper is that the CHT09 data analysis is seriously flawed even if one ignores the theoretically expected possibility that the core and envelope magnetic fields will have different directions.712 In this paper. we first present observational evidence independent of the CIITOO data. which shows considerable density structure in the four observed envelopes — thus suggesting fielc-streneth variations as well and then show that the new arguments of CLETIO are even more flawed than the original CLITOO analysis: they violate basic rules of logic and. scientific reasoning.," In this paper, we first present observational evidence independent of the CHT09 data, which shows considerable density structure in the four observed envelopes – thus suggesting field-strength variations as well – and then show that the new arguments of CHT10 are even more flawed than the original CHT09 analysis; they violate basic rules of logic and scientific reasoning."713 Our main objection to the CIEPOO. analvsis lies in the [act that they impose the overly restrictive of zero spread in the envelope D-field. values in each of the four clouds. without any evidence that the spread. is indeed zero. (," Our main objection to the CHT09 analysis lies in the fact that they impose the overly restrictive of zero spread in the envelope $B$ -field values in each of the four clouds, without any evidence that the spread is indeed zero. ("714Pheir own data show a preference. for an intrinsic spread approximately equal to the scatter induced by their (large) observational uncertainties see below.),Their own data show a preference for an intrinsic spread approximately equal to the scatter induced by their (large) observational uncertainties – see below.)715 lt is therefore relevant to ask whether there is any—imédependen observational evidence supporting or contradicting the assumption of no-spread in the envelope D-ield. values., It is therefore relevant to ask whether there is any observational evidence supporting or contradicting the assumption of no-spread in the envelope $B$ -field values.716 ‘To answer this question we plot intensity maps of Spitzer continuum emission (in 100. 70. and 24 yam. RGB colors) and C€O emission. (white contours). tracing the column density in the cores and their environments. in Fig.," To answer this question we plot intensity maps of Spitzer continuum emission (in 160, 70, and 24 $\mu$ m, RGB colors) and $^{13}$ CO emission (white contours), tracing the column density in the cores and their environments, in Fig."717 1. for L1448C'O and Dl. and in Fig.," \ref{perseus} for L1448CO and B1, and in Fig."718 2. [lor D217-2 and L1544. respectively (seo figure captions for CO and Spitzer data references).," \ref{taurus} for B217-2 and L1544, respectively (see figure captions for $^{13}$ CO and Spitzer data references)."719 On each map we overplot the Arecibo beam (black circle) used for the measurements of the core magnetic Ποια strengths. and the four GD'T beams (evan circles) used by CLEFO9 for the envelope magnetic-Teld measurements.," On each map we overplot the Arecibo beam (black circle) used for the measurements of the core magnetic field strengths, and the four GBT beams (cyan circles) used by CHT09 for the envelope magnetic-field measurements."720 lt is clear from these maps that the four GITE beams probe regions in the clouds! envelopes with very diverse morphologies and densities., It is clear from these maps that the four GBT beams probe regions in the clouds' envelopes with very diverse morphologies and densities.721 Phe expectation for the magnetic field. then is to exhibit a similar. diversity in morphology ancl magnitude., The expectation for the magnetic field then is to exhibit a similar diversity in morphology and magnitude.722 In. fact. one of the CLT authors (Crutcher 9010) has recently argued. that the magnetic field. (D) scales with density (p) as Dx(pE ," In fact, one of the CHT authors (Crutcher 2010) has recently argued that the magnetic field $B$ ) scales with density $\rho$ ) as $B \propto \rho^{2/3}$."723We disagree withthe exponent. of this relation. but we predicted long ago that a positive correlation between P and p (namely. Bxpl?) should indeed exist μμ isothermal. magnetically supported objects. (Mouschovias 1976: see also Fiedler Alousehovias 1993. Fig.," We disagree withthe exponent of this relation, but we predicted long ago that a positive correlation between $B$ and $\rho$ (namely, $B \propto \rho^{1/2}$ ) should indeed exist in self-gravitating, isothermal, magnetically supported objects (Mouschovias 1976; see also Fiedler Mouschovias 1993, Fig."724 9c. and review by Mouschovias 1996).," 9c, and review by Mouschovias 1996)."725 Llowever. if one adopts their more sensitive οp scaling. there is an even stronger reason to expect significant variation of the magnetic-field streneth in the cloud. envelopes. since there is considerable density structure there.," However, if one adopts their more sensitive $B - \rho$ scaling, there is an even stronger reason to expect significant variation of the magnetic-field strength in the cloud envelopes, since there is considerable density structure there."726 In other words.envelope.," In other words,."727 Our disagreement with CIIT on the treatment of their data is about and produce a single upper limit for I (the ratio of core and envelope mass-to-Llux ratio) in eachcloud., Our disagreement with CHT on the treatment of their data is about and produce a single upper limit for $R$ (the ratio of core and envelope mass-to-flux ratio) in eachcloud.728 An additional. important. disagreement stems from the [act that neither CIETO9 nor CIELO quotes upper limits on the quantity 2.," An additional, important disagreement stems from the fact that neither CHT09 nor CHT10 quotes upper limits on the quantity $R$."729 Upper limits are the only appropriate way do quote the information content in the CLETOO data. since in 15 out of 16 cases the mean magnetic-field strength in each cloud. envelope is consistent with zero (i... the measurements vielcleck nondetections).," Upper limits are the only appropriate way to quote the information content in the CHT09 data, since in 15 out of 16 cases the mean magnetic-field strength in each cloud envelope is consistent with zero (i.e., the measurements yielded nondetections)."730 In order to facilitate visual examination of the data the four envelope positions are combined. (which is where the CLEPO9/NLTOO. disagreement comes in play). we show in Fig.," In order to facilitate visual examination of the data the four envelope positions are combined (which is where the CHT09/MT09 disagreement comes in play), we show in Fig."731 3. the 8¢ upper limit derived for 2 (labeled 1. 2. 3. 4. corresponding to north. east. west. and south. respectively).," \ref{indones} the $3\sigma$ upper limit derived for $R$ (labeled 1, 2, 3, 4, corresponding to north, east, west, and south, respectively)."732 For L1544west there is à 37 detection for the envelope £-Lield.for which we can also derive a 30 measurement of Z7.," For L1544west there is a $3\sigma$ detection for the envelope $B$ -field,for which we can also derive a $3\sigma$ measurement of $R$ ."733 In this case. 2=3.5x 1.," In this case, $R=3.5 \pm 1$ ."734 Inall other cases. the individual 30 upper limits are consistent," Inall other cases, the individual $3\sigma$ upper limits are consistent"7350.0138AL.XRywp0152AL.. and 20.000IxXZr25.200 Ix. tthe GS per cent confidence ranges). we find a range of cooling ages between 11.3 and 50.1 million vears.,"$0.0138\,M_{\odot}\le R_{\rm WD}\le 0.0152\,M_{\odot}$, and $20,000~{\rm K}\le T_{\rm eff}\le 25,200$ K the 68 per cent confidence ranges), we find a range of cooling ages between 11.3 and 50.1 million years."736 The fact that the white dwarf appears to be relatively voung may put constraints on how the present-cay binary formed., The fact that the white dwarf appears to be relatively young may put constraints on how the present-day binary formed.737 For example. if the white dwarf formed first. then the second star would have had. S$50 million vears additional time to evolve into the sdB star. assuming the white cwarl has not somehow been heated since its formation.," For example, if the white dwarf formed first, then the second star would have had $\la 50$ million years additional time to evolve into the sdB star, assuming the white dwarf has not somehow been heated since its formation."738 For the second method to potentially reduce the parameter space found by the W-D fits we can exploit the fact that the mass of the sdB star is stronely correlated. with its racius (Oi is positively correlated with Q)., For the second method to potentially reduce the parameter space found by the W-D fits we can exploit the fact that the mass of the sdB star is strongly correlated with its radius $\Omega_{\rm sdB}$ is positively correlated with $Q$ ).739 Hence. the observed rotational velocity of the sdB star will be correlated. with its raclius.," Hence, the observed rotational velocity of the sdB star will be correlated with its radius."740 One must assume a rotational period for the sd star. however.," One must assume a rotational period for the sdB star, however."741 The usual assumption for close binary stars is that Po=Pan. ssvnchronous rotation. although his may not necessarily be the case (see below).," The usual assumption for close binary stars is that $P_{\rm rot}=P_{\rm orb}$, synchronous rotation, although this may not necessarily be the case (see below)."742 I£ the rotation is svnchronous. then Vssin?=SO km s+ when Αν=0.2.UM. and Viasing=125 km s when As=TÀAM.. where the values are computed. without using the white dwarf masseradius constraint.," If the rotation is synchronous, then $V_{\rm rot}\sin i=80$ km $^{-1}$ when $M_{\rm sdB}=0.2\,M_{\odot}$ and $V_{\rm rot}\sin i=125$ km $^{-1}$ when $M_{\rm sdB}=0.7\,M_{\odot}$ , where the values are computed without using the white dwarf mass-radius constraint."743 Thus a measurement of Visin?’ to better than =10 km would allow one o reduce the available parameter space for the mass of he sdl3 star., Thus a measurement of $V_{\rm rot}\sin i$ to better than $\approx 10$ km $^{-1}$ would allow one to reduce the available parameter space for the mass of the sdB star.744 The surface eravity computed. from the D fits is also correlated with the computed mass of the κα) star., The surface gravity computed from the W-D fits is also correlated with the computed mass of the sdB star.745 We find logg=5.46 (ees) when AL=0.2Al. and ogg=5.59- fees) when έως=0.7.M.. again where the values are computed without using the white cwarl mass-radius constraint.," We find $\log g=5.46$ (cgs) when $M_{\rm sdB}=0.2\,M_{\odot}$ and $\log g=5.59$ (cgs) when $M_{\rm sdB}=0.7\,M_{\odot}$, again where the values are computed without using the white dwarf mass-radius constraint."746 A spectroscopic measurement of logg is model dependent and requires high quality data (one slight advantage of using logg to constrain the range of sdB masses is that one does not need to assume a rotational period for the sdB star)., A spectroscopic measurement of $\log g$ is model dependent and requires high quality data (one slight advantage of using $\log g$ to constrain the range of sdB masses is that one does not need to assume a rotational period for the sdB star).747 Unfortunately. the spectroscopic measurement given by IXOW of logg=5440.1 does not provide a definitive answer.," Unfortunately, the spectroscopic measurement given by KOW of $\log g=5.4 \pm 0.1$ does not provide a definitive answer."748 The 20 range covers nearly all of the logg range of the W-D fits and Monte Carlo simulation., The $2\sigma$ range covers nearly all of the $\log g$ range of the W-D fits and Monte Carlo simulation.749 The more precisely defined. eclipse profiles also allow us to determine the orbital phase much more accurately., The more precisely defined eclipse profiles also allow us to determine the orbital phase much more accurately.750 The differential corrections routine of the W-D code can be usec to determinethe optimal phase (relative to some assume time) and. its probable error., The differential corrections routine of the W-D code can be used to determinethe optimal phase (relative to some assumed time) and its probable error.751 The heliocentric time of the superior conjunction of the sdD star is given in Table. 1., The heliocentric time of the superior conjunction of the sdB star is given in Table 1.752 The formal lo error on this time Is 4.9 seconds. compare with the error of 43 seconds on that measurement. given inW," The formal $1\sigma$ error on this time is 4.9 seconds, compared with the error of 43 seconds on that measurement given in."753OW? Our determination of the phase combined with IXOW's determination leads to the improved. measuremen of the orbital period given in Table 1., Our determination of the phase combined with KOW's determination leads to the improved measurement of the orbital period given in Table 1.754 The accuracy with which we can determine the orbita phase allows for the possibility of measuring the change in the binary. period., The accuracy with which we can determine the orbital phase allows for the possibility of measuring the change in the binary period.755 A possible cause ofa change in the orbita period. would be the loss of orbital angular momentum via the radiation of gravitational waves., A possible cause of a change in the orbital period would be the loss of orbital angular momentum via the radiation of gravitational waves.756 litter (1986) gives the period derivative due to gravitational wave radiation as where the masses are in solar masses and the period is in days., Ritter (1986) gives the period derivative due to gravitational wave radiation as where the masses are in solar masses and the period is in days.757 Using the masses computed: using the white chvarl mass-radius constraint and 7 given in Table 1 we lind Por=(3.1a)10 £7., Using the masses computed using the white dwarf mass-radius constraint and $P$ given in Table 1 we find $\dot{P}_{\rm GR}=-(3.1^{+0.4}_{-0.2})\times 10^{-13}$ .758" The phase dillerence Ao,ο which accumulates over N orbital eveles between the", The phase difference $\Delta\phi_{o-c}$ which accumulates over $N$ orbital cycles between the759'eview Low soar racliation pressure allects oryal parameters.,review how solar radiation pressure affects orbital parameters.760 This analysis roughly follows he treatinen eiven in Horáuyiaud.Burus((1991) for a particle in orbit around Jupiter. nt is generaized to account Lor tlie possibiity that the Stun may be located siguificautly above or below the ‘ineplane.," This analysis roughly follows the treatment given in \citet{HB91} for a particle in orbit around Jupiter, but is generalized to account for the possibility that the Sun may be located significantly above or below the ringplane."761 Also. we will 'strict ourselves to uearly circular orbits. hereby obtainingj sliipler expressions than tLose given by Hamilton(1993).," Also, we will restrict ourselves to nearly circular orbits, thereby obtaining simpler expressions than those given by \citet{Hamilton93}."762. Note that hroughout this analysis we assume the clyuamics of the particles is determined entirely by solar racdiatio κ.'e aud Saturn's gravity (0her uon-gravitationale forces such as plasiua drag are ueglec(ος)., Note that throughout this analysis we assume the dynamics of the particles is determined entirely by solar radiation pressure and Saturn's gravity (other non-gravitational forces such as plasma drag are neglected).763 We beegiu wil1 the standard perturbaion equations for the semi-major axis d. eccentricitv €. inelinaion 7. the longitude of periapse z aud the longitude of node £2 of a particle orbit (see e.g. Burus 1976).," \nocite{Burns76}764 We begin with the standard perturbation equations for the semi-major axis $a$ , eccentricity $e$, inclination $i$, the longitude of periapse $\varpi$ and the longitude of node $\Omega$ of a particle orbit (see e.g. Burns 1976)."765" Siuce we are iuterested in orbits with small eccentricities and inclinations. these expressions can be approximated as where n ls 1le paricle's wean motion. Fy;=GALingfar is approximately he fove ol Saturn's graviy OH i| particle with nass my (neglecting the eflects of Saturvs finite Oblateness). f is he paqlicle's true anoualy aud £,. Fy aud F. are the radial. azimutlal ancl normal (to the orbi plane in the cli‘ection of orbital angular momenttin) coiipoueils olf the perturbing force. respectively."," Since we are interested in orbits with small eccentricities and inclinations, these expressions can be approximated as where $n$ is the particle's mean motion, $F_{G}=GMm_g/a^2$ is approximately the force of Saturn's gravity on a particle with mass $m_g$ (neglecting the effects of Saturn's finite oblateness), $f$ is the particle's true anomaly and $F_r$, $F_t$ and $F_z$ are the radial, azimuthal and normal (to the orbit plane in the direction of orbital angular momentum) components of the perturbing force, respectively."766 Say the Suu is locaecd at au elevation augle B. above the rings aid a longittde A. in some inertial coordiuate system., Say the Sun is located at an elevation angle $B_\Sun$ above the rings and a longitude $\lambda_\Sun$ in some inertial coordinate system.767" Then the components of thesolar radiation pressure force FF, at a specilied loneitide A in the rine are given by:", Then the components of thesolar radiation pressure force $F_\Sun$ at a specified longitude $\lambda$ in the ring are given by:768eecnissiou offset from the FUV peak.,emission offset from the FUV peak.769 The couveutioual description of the star formation process begius with the (possibly trigeered) gravitational collapse of massive gas clouds iu the disk of a galaxy., The conventional description of the star formation process begins with the (possibly triggered) gravitational collapse of massive gas clouds in the disk of a galaxy.770" These ""super-clouds may be in the form of oorΠο. depending ou the plivsical conditions in the ISM EEhnneereen 1991 and references therein: Lhuesreen 1993)."," These “super-clouds” may be in the form of or, depending on the physical conditions in the ISM Elmegreen 1991 and references therein; Elmegreen 1993)."771 The denser parts of these clouds cool further (aud turni iuto uf they arent already in that forma) in giaut molecular clouds (GAIC's)}: stars then form aud evolve in these GAICs SShu et 1993)., The denser parts of these clouds cool further (and turn into if they aren't already in that form) in giant molecular clouds (GMCs); stars then form and evolve in these GMCs Shu et 1993).772" In this picture. a simall fraction of the eas is coustuned in the star formation process, aud a larger fraction may be ionized aucd/or evacuated by stellar winds and superunovae."," In this picture, a small fraction of the gas is consumed in the star formation process, and a larger fraction may be ionized and/or evacuated by stellar winds and supernovae."773 Some photodissociation of into hhas also been observed on the surfaces of specific GAICs in our Galaxy DDlitz 1993. section 6: Andersson λαο 1993: Kuchar Dania 1993) and ουπα emP?deusitiespcof.order41«+1023Di hhave been reported: but the photodissociation process is usually uot thought of as being responsible for producing a substautial fraction of the GalacticIT1.," Some photodissociation of into has also been observed on the surfaces of specific GMCs in our Galaxy Blitz 1993, section 6; Andersson Wannier 1993; Kuchar Bania 1993) and column densities of order } have been reported; but the photodissociation process is usually not thought of as being responsible for producing a substantial fraction of the Galactic."774. However. our model for the morphology of the UV. Ue. and Hin ADSI leads to the conclusion that a laree fraction. and possibly all. of the Hn the inuer spiral aris of that galaxy is a photodissociation product.," However, our model for the morphology of the UV, $\alpha$, and in M81 leads to the conclusion that a large fraction, and possibly all, of the in the inner spiral arms of that galaxy is a photodissociation product."775" TheILL. Ho. and PUY structures observed in M81 can be understood if ""οσον, like those proposed by Norman Ikeuchi (1989). are conuuon in the spiral aris of this galaxy."," The, $\alpha$, and FUV structures observed in M81 can be understood if “chimneys”, like those proposed by Norman Ikeuchi (1989), are common in the spiral arms of this galaxy."776 In this picture. holes 100 pc in size are blown out of the galactic disk bv concentrations of ligh-mass star formation activity. simular to that discovered above the Galactic ireegion ICls05 bv Nonuaudeau. Tavlor. Dewduey (1996).," In this picture, holes 100 pc in size are blown out of the galactic disk by concentrations of high-mass star formation activity, similar to that discovered above the Galactic region IC1805 by Normandeau, Taylor, Dewdney (1996)."777 From our vautage point above the disk of M81. such concentrations of star-forming activity are seen essentially free of obscuration. since the lot stars have evacuated or destroved most of the iuterveuiug dust in M81.," From our vantage point above the disk of M81, such concentrations of star-forming activity are seen essentially free of obscuration, since the hot stars have evacuated or destroyed most of the intervening dust in M81."778 The eas in MSIE is largely in the form of at least iu the main part of the disk: ls produced extensively by UV photons from the voung stars muipingiug on the mner surfaces of these structures., The gas in M81 is largely in the form of at least in the main part of the disk; is produced extensively by UV photons from the young stars impinging on the inner surfaces of these structures.779" The detailed structure of such an Haver may reseible the iunorphologv of the Galactic star-forming region C216-2.5, which has been explained as a PDR by Williams Maddalena (1996)."," The detailed structure of such an layer may resemble the morphology of the Galactic star-forming region G216-2.5, which has been explained as a PDR by Williams Maddalena (1996)."780 This region shows a laver of zz50 pe thick spread over 300 pe on the outer surface of a molecular cloud: the οσο density is of order cmThe0., This region shows a layer of $\approx 50$ pc thick spread over 300 pc on the outer surface of a molecular cloud; the column density is of order } .78121031 2... cause of the photodissociation is thought to be two voung stars located ~50 pe from the laver., The cause of the photodissociation is thought to be two young stars located $\sim 50$ pc from the layer.782 Tu Fig.7 we show a simplified cross-sectional sketeli of the various morphologies which may arise m such a picture., In \ref{cartoon} we show a simplified cross-sectional sketch of the various morphologies which may arise in such a picture.783 The observer views the galaxy disk from above., The observer views the galaxy disk from above.784 Case I illustrates the 161505 ecometry. which forms the basis for our model.," Case I illustrates the IC1805 geometry, which forms the basis for our model."785 A voune star cluster has formed just above the mid-plane in a dense laver of molecular gas., A young star cluster has formed just above the mid-plane in a dense layer of molecular gas.786 The Haver is shown in the sketch as a smooth medium but will have structure., The layer is shown in the sketch as a smooth medium but will have structure.787 If O stars are present. the accompanving Ho will be roughly coincident with the UV.," If O stars are present, the accompanying $\alpha$ will be roughly coincident with the UV."788 Since the star-formation activity has opened a chimney up into the halo. the observer sees the optical and UV emission with little obscuration. even though he laver underneath the bottom of the chinmey will contain dust aud may be opaque.," Since the star-formation activity has opened a chimney up into the halo, the observer sees the optical and UV emission with little obscuration, even though the layer underneath the bottom of the chimney will contain dust and may be opaque."789 TheILL. produced youn the iy pliotodissociation. is spread over the πιο surface of the chimney aud will appear to the observer to f enhanced at the edges of the region. owiug to he longer lines of sight there.," The, produced from the by photodissociation, is spread over the inner surface of the chimney and will appear to the observer to be enhanced at the edges of the region, owing to the longer lines of sight there."790 Case IT is the same ecolctry as Case [but since the chimney has blown out on the far side of a more-or-less opaque disk of eas. the observer may measure nothing more than a patch of eecnission.," Case II is the same geometry as Case I, but since the chimney has blown out on the far side of a more-or-less opaque disk of gas, the observer may measure nothing more than a patch of emission."791 Case III in Fig.7 shows a particularly cucrectic star-forming region located very close to the iuid-plane of the galaxy., Case III in \ref{cartoon} shows a particularly energetic star-forming region located very close to the mid-plane of the galaxy.792 Iu this case. chimneys may blow out on both sices of the disk. leaving a hole right through the galaxy.," In this case, chimneys may blow out on both sides of the disk, leaving a hole right through the galaxy."793" The photodissociated wwill appear to the observer as an ""Ueocoon wrapped around the bouudary of the UV-uitting region. as seen nn Field 3."," The photodissociated will appear to the observer as an “cocoon” wrapped around the boundary of the UV-emitting region, as seen in Field 3."794 Case IV is a one-sided version of Case TT. formed ina nou-unitforui part of the," Case IV is a one-sided version of Case III, formed in a non-uniform part of the"795fraction of the intrinsic continuum into the line-of-sight.,fraction of the intrinsic continuum into the line-of-sight.796 If this scattered continuum is the dominant observec continuum (as would be the case if the direct line-of-sight is obscured by Compton-thick matter). then the EW of the emission lines from ionized Fe could be large. of the order of hundreds of eV or more.," If this scattered continuum is the dominant observed continuum (as would be the case if the direct line-of-sight is obscured by Compton-thick matter), then the EW of the emission lines from ionized Fe could be large, of the order of hundreds of eV or more."797 The details depend on several factors. principally the ionization parameter and column density of the line-emitting region. as well as the shape of the tonizing continuum (see for example. detailed calculations of the EW of the Fe and Fe emission lines in Bianchi&Matt 2002)).," The details depend on several factors, principally the ionization parameter and column density of the line-emitting region, as well as the shape of the ionizing continuum (see for example, detailed calculations of the EW of the Fe and Fe emission lines in \citealt{Bianchi2002}) )."798 Suppose that the Compton-thick reprocessor subtends a solid angle AQ/4z at the X-ray source and that the system is observed along a line-of-sight that does not give a direct view of the X-ray source and that has the greatest column density. Nj. (this maximizes the Fe K line EW).," Suppose that the Compton-thick reprocessor subtends a solid angle $\Delta \Omega/4\pi$ at the X-ray source and that the system is observed along a line-of-sight that does not give a direct view of the X-ray source and that has the greatest column density, $N_{H}$, (this maximizes the Fe K line EW)."799 Further suppose that the space between the reprocessor is filled by à warm. optically-thin scattering zone with Thomson depth rj; (with zi;<< 1D. subtending a solid angle |—(AQ/47) at the X-ray source.," Further suppose that the space between the reprocessor is filled by a warm, optically-thin scattering zone with Thomson depth $\tau_{\rm thin}$ (with $\tau_{\rm thin} <<1 $ ), subtending a solid angle $1-(\Delta \Omega/4\pi)$ at the X-ray source."800 Then. a fraction f=tyin[1]—(AO/47)] of the intrinsic X-ray continuum is scattered into the observer's line-of-sight and will reduce the EW of the Fe K line if it dominates over the zeroth-order continuum that is observed directly through the Compton-thick absorber.," Then, a fraction $f\equiv \tau_{\rm thin}[1-(\Delta \Omega/4\pi)]$ of the intrinsic X-ray continuum is scattered into the observer's line-of-sight and will reduce the EW of the Fe K line if it dominates over the zeroth-order continuum that is observed directly through the Compton-thick absorber."801" Specifically. for column densities greater than a few x107em"". even a small value of f can significantly reduce the EW of the FeK line."," Specifically, for column densities greater than a few $\times80210^{24} \ \rm cm^{-2}$, even a small value of $f$ can significantly reduce the EW of the FeK line."803 For example. for a column density of 107?cem. a scattering fraction of f>0.01 will reduce the EW of an Fe K line by more than an order of magnitude. so that an EW of 1 keV would be reduced to less than ~100 eV. and it could render the line undetectable (see Ghisellinietal. 1994)).," For example, for a column density of $10^{25} \ \rm cm^{-2}$, a scattering fraction of $f>0.01$ will reduce the EW of an Fe K line by more than an order of magnitude, so that an EW of 1 keV would be reduced to less than $\sim 100$ eV, and it could render the line undetectable (see \citealt{Ghisellini}) )."804 Furthermore. the intrinsic EW (re. prior to the dilution effect) of Fe K lines depends on several factors not only the column density of the absorber but also the geometry of the absorber (e.g. the half opening angle of the putative torus; see Ghisellinietal.1994;Ikeda2009;Matt 1996)).," Furthermore, the intrinsic EW (i.e. prior to the dilution effect) of Fe K lines depends on several factors not only the column density of the absorber but also the geometry of the absorber (e.g. the half opening angle of the putative torus; see \citealt{Ghisellini,Ikeda,Matt}) )."805 For example for an half opening angle of 30* and our estimate of the column density of the neutral absorber the intrinsic EW of the 6.4 keV Fe line can span the range from | to 4 keV (Ghisellinietal., For example for an half opening angle of $30^{\circ}$ and our estimate of the column density of the neutral absorber the intrinsic EW of the 6.4 keV Fe line can span the range from 1 to 4 keV \citep{Ghisellini}.806.1994).. The inferred column density of the Compton-thick reprocessor implies. that a scattering. fraction of only ~0.1% in the optically-thin zone is required to begin to dilute the 6.4 keV Fe K line and a scattering fraction of a few percent is sufficient to reduce the EW of the line well below 100 eV. consistent with the upper limit of the EW and the ~2% scattering fraction (measured with respect to the de-absorbed primary power law component) as measured with the Suzaku data.," The inferred column density of the Compton-thick reprocessor implies that a scattering fraction of only $\sim 0.1\%$ in the optically-thin zone is required to begin to dilute the 6.4 keV Fe K line and a scattering fraction of a few percent is sufficient to reduce the EW of the line well below 100 eV, consistent with the upper limit of the EW and the $\sim 2$ scattering fraction (measured with respect to the de-absorbed primary power law component) as measured with the Suzaku data."807 A second possible geometry is that we have a direct view of the inner surface of the Compton-thick reprocessor. but the outer part of this reprocessor is 1onized.," A second possible geometry is that we have a direct view of the inner surface of the Compton-thick reprocessor, but the outer part of this reprocessor is ionized."808 In this case. if the remaining part of the reprocessor is Compton-thin. the EW of the 6.4 keV Fe Ka line will be reduced and the emission detected below 10 keV ts the reflected emission from this inner ionized surface of the torus which will also produce a strong 6.7 keV line.," In this case, if the remaining part of the reprocessor is Compton-thin, the EW of the 6.4 keV Fe $\alpha$ line will be reduced and the emission detected below 10 keV is the reflected emission from this inner ionized surface of the torus which will also produce a strong 6.7 keV line."809 Thus. we see that the lack of a large EW neutral Fe K in IRAS 19254—7245 1s not unexpected.," Thus, we see that the lack of a large EW neutral Fe K in IRAS $-$ 7245 is not unexpected."810 The spectrum below 10 keV is then dominated by this optically-thin scattered continuum and the dominance of the emission line from ionized Fe is consistent with this It is worth noting that this ts not a unique case of a detection of a strong 6.7 keV line in a luminous infrared galaxy., The spectrum below 10 keV is then dominated by this optically-thin scattered continuum and the dominance of the emission line from ionized Fe is consistent with this It is worth noting that this is not a unique case of a detection of a strong 6.7 keV line in a luminous infrared galaxy.811 Other examples are Arp299 (Balloetal..2004).. Arp220 (Iwasawaal..2005:Tenget2009).. and IRAS 00182-7112 (Nandra&Iwasawa.2007).," Other examples are Arp299 \citep{ballo04}, Arp220 \citep{iwasawa05,Teng}, and IRAS 00182-7112 \citep{Nandra07}."812. For all these sources although the optical spectra show no clear signature of AGN activity. their X-ray emission and the 6.7 keV line can be explained with the presence of an AGN and an tonized reflector as in the case of19254-7245.," For all these sources although the optical spectra show no clear signature of AGN activity, their X-ray emission and the 6.7 keV line can be explained with the presence of an AGN and an ionized reflector as in the case of."813. However. while in the case of Arp 299 and Arp 220. the X-ray luminosity is not indicative that the major contributor to the bolometric luminosity is a high-luminosity AGN. in the case of IRAS 00182-7112 the X-ray luminosity is too large to be accounted for by the strong starforming activity (LQ.-10)>IOUereg so!) as for19254-7245.," However, while in the case of Arp 299 and Arp 220, the X-ray luminosity is not indicative that the major contributor to the bolometric luminosity is a high-luminosity AGN, in the case of IRAS 00182-7112 the X-ray luminosity is too large to be accounted for by the strong starforming activity $(2-10)> 10^{44}$ ) as for."814. For all these sources. the presence of a strong tonized Fe line. with little or no 6.4 keV line. could be reconciled with the picture of a heavily obscured AGN assuming that we do not have a direct view of the reflected continuum. because it is is diluted by the scattering from the tonized matter. that is associated with the production of the line. or if the surface of the putative Compton-thick reprocessor ts highly Overall to account for the X-ray emission of above 2 keV. we need two absorbing/reflecting media: one neutral and Compton-Thick and probably seen in transmission. and one ionized and probably seen in reflection.," For all these sources, the presence of a strong ionized Fe line, with little or no 6.4 keV line, could be reconciled with the picture of a heavily obscured AGN assuming that we do not have a direct view of the reflected continuum, because it is is diluted by the scattering from the ionized matter, that is associated with the production of the line, or if the surface of the putative Compton-thick reprocessor is highly Overall to account for the X-ray emission of above 2 keV, we need two absorbing/reflecting media: one neutral and Compton-Thick and probably seen in transmission, and one ionized and probably seen in reflection."815 This latter is responsible for the flat X-ray spectrum emerging below 10 keV. for the He-like Fe K line and probably for the dilution of the 6.4 keV Fe line produced in the neutral Compton-Thick Despite the various possible models for the 0.9--10 keV emission. we always need a neutral high column density absorber to account for the emission emerging above 10 keV. Once we have corrected for the amount of absorption. the intrinsic 2-10 keV luminosity of the primary AG component is ~107 erg s7!.," This latter is responsible for the flat X-ray spectrum emerging below 10 keV, for the He-like Fe K line and probably for the dilution of the 6.4 keV Fe line produced in the neutral Compton-Thick Despite the various possible models for the 0.5–10 keV emission, we always need a neutral high column density absorber to account for the emission emerging above 10 keV. Once we have corrected for the amount of absorption, the intrinsic 2–10 keV luminosity of the primary AGN component is $\sim10^{44}$ erg $^{-1}$."816 This high column density absorber may also be the one responsible for the deep hydrocarbon absorption detected in the L-band spectrum (Risalitietal.2003) at 3.4yun., This high column density absorber may also be the one responsible for the deep hydrocarbon absorption detected in the L-band spectrum \citep{Risaliti2003} at $\mu m$.817 Finally. independently from the assumed model for the 2-35 keV emission (re. tonized reflection or scattered power law component). we always require the soft thermal component to account for the 0.5-2 keV emission.," Finally, independently from the assumed model for the 2–35 keV emission (i.e. ionized reflection or scattered power law component), we always require the soft thermal component to account for the 0.5–2 keV emission."818 As already found for other ULIRGs this thermal component has a temperature AT~0.7 keV: the luminosity of this component is L(0.85-2keV)~4x104! s7!.., As already found for other ULIRGs this thermal component has a temperature $kT\sim 0.7$ keV; the luminosity of this component is $L(0.5-2 \;\mathrm{keV})\sim 4\times 10^{41}$ .819 Although. we cannot exclude a possible contribution from the ionized reflector. the measured soft X-ray luminosity is in agreement with that expected from," Although, we cannot exclude a possible contribution from the ionized reflector, the measured soft X-ray luminosity is in agreement with that expected from"82011997. Scott οἱ 22000b).,"1997, Scott et 2000b)."821 The results. summarized in Paper II of this series (Scott el 22000b). are in general agreement wilh (he predictions of models of the UV background which integrate the contribution Irom known population of quasars aud imclude reprocessing ellects in an inhomogeneous intergalactie medium (IHaardt Aladan 1996. hereafter HAI9G. Farclal οἱ 11998).," The results, summarized in Paper II of this series (Scott et 2000b), are in general agreement with the predictions of models of the UV background which integrate the contribution from known population of quasars and include reprocessing effects in an inhomogeneous intergalactic medium (Haardt Madau 1996, hereafter HM96, Fardal et 1998)."822 In Paper IH. the mean intensity of the ionizing background was found to be 7.0tix107? eres bem 7 F tat ze3.," In Paper II, the mean intensity of the ionizing background was found to be $7.0^{+3.4}_{-4.4} \times 10^{-22}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ at $z \sim 3$."823 The decline of the quasar space density from ze2 to the present is expected to drive a corresponding decline in the intensity of the UV background., The decline of the quasar space density from $z \sim 2$ to the present is expected to drive a corresponding decline in the intensity of the UV background.824 Iulkarni Fall (1993. hereafter KF93) measured η)6x107! eres 1 ? F1 tat το0.5 [rom a subset of the now complete IST Quasar Absorption Line Kev Project sample presented by ασ et ((1993).," Kulkarni Fall (1993, hereafter KF93) measured $J(\nu_{0}) \sim 6825\times 10^{-24}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ at $z \sim 0.5$ from a subset of the now complete HST Quasar Absorption Line Key Project sample presented by Bahcall et (1993)."826 Much of (his previous work has relied upon the technique lor measuring (1) outlined by BDO., Much of this previous work has relied upon the technique for measuring $J(\nu_{0})$ outlined by BDO.827 This technique requires the entire sample of absorption lines to be binned according to the ratio of the quasar [lux at the physical position of the absorber to the backeround flux., This technique requires the entire sample of absorption lines to be binned according to the ratio of the quasar flux at the physical position of the absorber to the background flux.828 This is done lor several initial guesses of the background intensity: and the value that gives the lowest 47 between the binned data and the ionization model is chosen as the best fit (79)., This is done for several initial guesses of the background intensity; and the value that gives the lowest $\chi^{2}$ between the binned data and the ionization model is chosen as the best fit $J(\nu_{0})$.829 ILowever. this is not the optimal technique to use at low redshift where absorption line densities are low.," However, this is not the optimal technique to use at low redshift where absorption line densities are low."830 IXF93 developed a maximun likelihood technique to address (his issue and used it in their measurement ο. έν) al 2~0.5., KF93 developed a maximum likelihood technique to address this issue and used it in their measurement of $J(\nu_{0})$ at $z \sim 0.5$.831 However. their measurement was based upon a sample of only 13 QSOs and less (han 100 lines. aud has correspondingly large error bars.," However, their measurement was based upon a sample of only 13 QSOs and less than 100 lines, and has correspondingly large error bars."832 In addition. (he value these authors find is lower than the predictions of the models ol Haardt Madau (1996). though consistent within the uneertainties. as shown in Figure 13 of Paper 11 and in Figure 11. of this paper.," In addition, the value these authors find is lower than the predictions of the models of Haardt Madau (1996), though consistent within the uncertainties, as shown in Figure 13 of Paper II and in Figure \ref{fig:lowzcomp} of this paper."833 Given the importance of the value of the LI ionization rate to the hyvdrodvnanmical evolution of the low redshlift universe. performing this measurement with a much larger line sample is worthwhile.," Given the importance of the value of the HI ionization rate to the hydrodynamical evolution of the low redshift universe, performing this measurement with a much larger line sample is worthwhile."834 The low redshift hvdrodsnamie simulations of Theuns et ((1998) ancl Dave οἱ ((1999) indicate that the evolution of the ionizing background is the primary driver behind the change of character of the Ly-a forest [rom hieh redshift to low redshilt. specifically. the break in the number distribution of Lv-a lines at z=1.7 (Morris et 11991. Baheall et 11991. Wevinann et 11998).," The low redshift hydrodynamic simulations of Theuns et (1998) and Davé et (1999) indicate that the evolution of the ionizing background is the primary driver behind the change of character of the $\alpha$ forest from high redshift to low redshift, specifically, the break in the number distribution of $\alpha$ lines at $z=1.7$ (Morris et 1991, Bahcall et 1991, Weymann et 1998)."835 The erowth of structure pulling eas from low density regions into high clensity regions also contributes to this and other attributes of the evolution of the Lv-o [orest., The growth of structure pulling gas from low density regions into high density regions also contributes to this and other attributes of the evolution of the $\alpha$ forest.836 Shull et ((1999) estimate the local ionizing background including contributions to the background from starburst galaxies as well as Sevler(s and QSOs., Shull et (1999) estimate the local ionizing background including contributions to the background from starburst galaxies as well as Seyferts and QSOs.837 Their models include a treatment of (he opacity of the low redshift Ly-a forest using information drawn from recent observational work (Wevimann et 11998. Penton et 22000b).," Their models include a treatment of the opacity of the low redshift $\alpha$ forest using information drawn from recent observational work (Weymann et 1998, Penton et 2000b)."838 Thev find (hat starbursts and AGN could contribute approximately equally to the ionizing background at low recshilt.," They find that starbursts and AGN could contribute approximately equally to the ionizing background at low redshift,"839for the PSPC sources in 1H detectedοπής of tle D25 ellipse is given in Table Al iu the appeudix (Appcucix A:)).,for the PSPC sources in 1 detected of the D25 ellipse is given in Table \ref{table_PsrcA} in the appendix \ref{sec_appe}) ).840 Ax discussed carlicr. we expected the total PSPC Observeution to be far more sensa.ive than the total URI Observeution.," As discussed earlier, we expected the total PSPC observation to be far more sensitive than the total HRI observation."841 Nevertheless we rediced the entire TRI data collected at three epochs after the PSPC observations (sec Table 1)) to check for time variailitv of sources. resolve confusc sotrees aud to obtain improved source positions.," Nevertheless we reduced the entire HRI data collected at three epochs after the PSPC observations (see Table \ref{table_obse}) ) to check for time variability of sources, resolve confused sources and to obtain improved source positions."842 For the analvsis we screened t16 observations for good time 1tervals longer thin sx: we nnde no further selectiums on low background tines (which would have reduceL the accepted time by 32:4)) as we were mainly iuteresed in poiut-like sources and for this purpose. we WOYO Sil proton lanited.," For the analysis we screened the observations for good time intervals longer than s; we made no further selections on low background times (which would have reduced the accepted time by ) as we were mainly interested in point-like sources and for this purpose, we were still photon limited."843 Before mereing the «atasets for source detection. we checked the attitude solution of the ineividial observations using three relatively bright sources detected in all 3 observations (IE. IT13. axd II16 of 33).," Before merging the datasets for source detection, we checked the attitude solution of the individual observations using three relatively bright sources detected in all 3 observations (H3, H13, and H16 of 3)."844" The positions of these sotrces in the dividual poiutiues aeree to better than c"".. and we therefore did not correct the attitude of the pointines before mereiug the data."," The positions of these sources in the individual pointings agree to better than $\pm 1$, and we therefore did not correct the attitude of the pointings before merging the data."845 As with the PSPC. soi rcotetection and position στοάiion was then perforna over the full feld of view ou the merecd dataset wih the EXSAS local detect. nap detect. and maxim111 ikclihood algorithuus (Zinineraun 11991). using muages of pixel sizeο restricted to IRI raw chainels 1.8 to reduce the cetector backerouud.," As with the PSPC, source detection and position determination was then performed over the full field of view on the merged dataset with the EXSAS local detect, map detect, and maximum likelihood algorithms (Zimmermann 1994), using images of pixel size, restricted to HRI raw channels 1–8 to reduce the detector background."846 We acceptec sources with a likelibood =8 as URI detections., We accepted sources with a likelihood $\ge 8$ as HRI detections.847 The resulting source list was then shifted with respect to the sav ΠΡ103. as done to nuprove the PSPC field positioning.," The resulting source list was then shifted with respect to the star HD5403, as done to improve the PSPC field positioning."848" Fiewe33 shows a full. kks IIRI contour image (channels S) of the central ~29"" οπή ireeion equivalent iu area to 11 2)."," 3 shows a full, ks HRI contour image (channels $-$ 8) of the central $\sim$ $\times$ region equivalent in area to 1 2)."849 Within the area covered by 3:3. Ls HRI sources are detected (with a likelihood £ 28). their source numbers marked in the figure. LO of which lie within the optical disk of NGC 300 (as again indicated bv the D25 cllipse).," Within the area covered by 3, 18 HRI sources are detected (with a likelihood $L$ $>$ 8), their source numbers marked in the figure, 10 of which lie within the optical disk of NGC 300 (as again indicated by the D25 ellipse)."850Mathur. S.. Wilkes. D..,"Mathur, S., Wilkes, B.,"851The author is thankful to Dr. Lue Dame and Dr. Javaraiah for (he useful discussions.,The author is thankful to Dr. Luc Dame and Dr. Javaraiah for the useful discussions.852in which he identifios between 7T and 119 YSCs that he calls Super Star Clusters (SSCs). and that we are currently analyzing in the way described above.,"in which he identifies between 7 and 149 YSCs that he calls Super Star Clusters (SSCs), and that we are currently analyzing in the way described above."853 All of them are compact with radii iu the range of 3 to 10 pe., All of them are compact with radii in the range of 3 to 10 pc.854 First of all. Ud like to caution the notion SSC. since it ouly refers toBaninositg aud uot to1055s.," First of all, I'd like to caution the notion SSC, since it only refers to and not to."855 The term SSC eoes back to van deu Dergh (7?) who referred to SCs much brighter than the brightest open clusters known in the Milkv. Wavy by that time., The term SSC goes back to van den Bergh \citep{vandenBergh71} who referred to SCs much brighter than the brightest open clusters known in the Milky Way by that time.856 Meamwhile. however. we know that other galaxies (e.g. the LAIC) cau have much richer SCs than our Milkv Wav aud we have detailed evolutionary svuthesis models that show how strouely SCs fade. in particular durius their vounecst stages.," Meanwhile, however, we know that other galaxies (e.g. the LMC) can have much richer SCs than our Milky Way and we have detailed evolutionary synthesis models that show how strongly SCs fade, in particular during their youngest stages."857 Depending on metallicity. a SC fades by ~15 mae in V dunug the first few hundred Myr alone through stellar evolution effects. ic. with the stellardvuamical mass loss not vet iucluded.," Depending on metallicity, a SC fades by $\sim 4 - 5$ mag in V during the first few hundred Myr – alone through stellar evolution effects, i.e. with the stellar-dynamical mass loss not yet included."858 A very buuinous SC therefore need not necessarily be extremely massive., A very luminous SC therefore need not necessarily be extremely massive.859 Even at relatively modest masses. SCs can bebright aud look like SSCs as long as they areyoung.," Even at relatively modest masses, SCs can be and look like SSCs as long as they are."860 We therefore strougly suggest to refer tomesses rather than tohaninositics for ¥SC's., We therefore strongly suggest to refer to rather than to for YSCs.861" The two brightest of LOUs. YSCs lave spectroscopic nasses of [107 aud 2-109NL, (2)..", The two brightest of L04's YSCs have spectroscopic masses of $4 \cdot 10^5$ and ${\rm 2 \cdot 10^6 ~ M_{\odot}}$ \citep{Larsen+04}.862 Being so few. these wo could. however. be singular outstanding SCs that ornmed by some local extreme compression of 1iolecular eas. e.g. from a supereiant molecular cloud compressed )etxyeen expanding shells.," Being so few, these two could, however, be singular outstanding SCs that formed by some local extreme compression of molecular gas, e.g. from a supergiant molecular cloud compressed between expanding shells."863 Our preliminary analvsis of the 321 YSC's in those six of LOUs ealaxies that have =30 YSCs cach. revealed hat a significant uuniber of ~70 YSCs with ages =KO Myr have masses 107ADL.," Our preliminary analysis of the 324 YSCs in those six of L04's galaxies that have $\geq 30$ YSCs each, revealed that a significant number of $\sim 70$ YSCs with ages $\geq 50$ Myr have masses ${\rm \geq 10^5~M_{\odot}}$."864 We have chosen a generous lower age limit of 50 Myr to be sure that the YSCs have already survived the most dangerous phase in their lives. the infant mortality phase after the first SNe have expelled the left-over eas aud the subsequeut dynamical rearrangcment to the change in the potential (cf.," We have chosen a generous lower age limit of 50 Myr to be sure that the YSCs have already survived the most dangerous phase in their lives, the infant mortality phase after the first SNe have expelled the left-over gas and the subsequent dynamical rearrangement to the change in the potential (cf."865 Lamers. vole).," Lamers, )."866 Aud we concentrate on YSC's with masses DLO?AL. xinee those lave fam survival chances for the forthcoming Crs according to prescut knowledge (7) and they have masses in the range of GCs. hence merit to be called voung GCs.," And we concentrate on YSCs with masses ${\rm \geq 10^5~M_{\odot}}$ since those have fair survival chances for the forthcoming Gyrs according to present knowledge \citep{BoutloukosLamers03} and they have masses in the range of GCs, hence merit to be called young GCs."867 This result that apparently uuidisturbed aud not currently starbursting She...Sd-type spirals form SCs which have all the properties of voung CCS ijs surprising and presents a challenge to our current wuderstanding of SC and GC formation aud evolution., This result that apparently undisturbed and not currently starbursting Sbc...Sd-type spirals form SCs which have all the properties of young GCs is surprising and presents a challenge to our current understanding of SC and GC formation and evolution.868 ? arenes8 that these spirals are not currently im any particularly active state of SE. they look uudisturbed aud feature nice disks witli regular spiral structure.," \cite*{Larsen04} argues that these spirals are not currently in any particularly active state of SF, they look undisturbed and feature nice disks with regular spiral structure."869 The age distributions that we obtain for their YSCs support this argunenut., The age distributions that we obtain for their YSCs support this argument.870 But how cau these undisturbed spirals afford the ligh SE effcieucies that current theories for GC formation require?, But how can these undisturbed spirals afford the high SF efficiencies that current theories for GC formation require?871 And where are the successors of previous eenerations of this type of SCs?, And where are the successors of previous generations of this type of SCs?872 Do normal actively star-forming spirals feature continuous age distributions amoue thei CC systems?, Do normal actively star-forming spirals feature continuous age distributions among their GC systems?873 Is our Milky Wav (and M31) special in this respect?, Is our Milky Way (and M31) special in this respect?874 Can intermediateage CC's have escaped our detection so far?, Can intermediate-age GCs have escaped our detection so far?875 Or are these actively star-forming galaxies at the same time particularly hostile to their YSCs aud destroy them bevond our current estimates2, Or are these actively star-forming galaxies at the same time particularly hostile to their YSCs and destroy them beyond our current estimates?876 We discuss this in more depth in Fritze et al..prep.," We discuss this in more depth in Fritze et al.,."877" Iu any case, our results require a careful vecousideratiouli of currently accepted coucepts of SC formation. evolution and destruction."," In any case, our results require a careful reconsideration of currently accepted concepts of SC formation, evolution, and destruction."878thermal time scale readjustment of the donor star to changes in the size of its Roche lobe.,thermal time scale readjustment of the donor star to changes in the size of its Roche lobe.879 Mass transfer typically starts when the donor star has a mass that is within a few tenths of a solar mass of the NS. and is also slightly evolved (with core mass ~0.1—0.25).," Mass transfer typically starts when the donor star has a mass that is within a few tenths of a solar mass of the NS, and is also slightly evolved (with core mass $\sim 0.1-0.25$ )."880 Because the donor is more massive than 0.8M... this type of system can be found only in younger GCs and/or in GCs with such a high rate of stellar interactions that at least one NS can be expected to have a blue straggler donor.," Because the donor is more massive than $0.8\, M_\odot,$ this type of system can be found only in younger GCs and/or in GCs with such a high rate of stellar interactions that at least one NS can be expected to have a blue straggler donor."881 One proposed test of this model was to check if bright X-ray sources in the GCSs of external galaxies are likely to be found in young and/or optically bright GCs., One proposed test of this model was to check if bright X-ray sources in the GCSs of external galaxies are likely to be found in young and/or optically bright GCs.882 This is exactly what we find in the Sombrero., This is exactly what we find in the Sombrero.883 Thus. although the model is not proved. M104 does provide some indirect evidence in its favor.," Thus, although the model is not proved, M104 does provide some indirect evidence in its favor."884 The correlation between optically bright GCs and high-luminosity X-ray sources ts intriguing., The correlation between optically bright GCs and high-luminosity X-ray sources is intriguing.885 It is likely that the large optical flux is indicative of a massive GC. in which stellar interactions producing X-ray binaries may be more frequent.," It is likely that the large optical flux is indicative of a massive GC, in which stellar interactions producing X-ray binaries may be more frequent."886 More frequent interactions may lead to more X-ray sources per cluster., More frequent interactions may lead to more X-ray sources per cluster.887 There are two possible checks for multiplicity., There are two possible checks for multiplicity.888 GCs that do not house multiples. or that include only a small number (~ 2—3) of individual components. could exhibit significant X-ray variability (by factors of 2—3) over the same time scales on which X-ray binaries are typically variable.," GCs that do not house multiples, or that include only a small number $\sim 2-3$ ) of individual components, could exhibit significant X-ray variability (by factors of $2-3$ ) over the same time scales on which X-ray binaries are typically variable."889 Thus. over time scales of days. weeks. or months. we could expect to find evidence of variability.," Thus, over time scales of days, weeks, or months, we could expect to find evidence of variability."890" In contrast. GCs that have a high level of X-ray source multiplicity-i.e.. in which the total flux is a composite consisting of comparable contributions from more than a few sources-should be almost constant. since it 15 unlikely that many of the independent components will have {, declining or increasing simultaneously."," In contrast, GCs that have a high level of X-ray source multiplicity–i.e., in which the total flux is a composite consisting of comparable contributions from more than a few sources–should be almost constant, since it is unlikely that many of the independent components will have $L$ declining or increasing simultaneously."891 At the distance of MIO4. one aresecond corresponds to roughly 43 pe. larger than the radii of typical GCs.," At the distance of M104, one arcsecond corresponds to roughly $43$ pc, larger than the radii of typical GCs."892 Most of the GC X-ray sources in our sample are located away from the aim point. in regions where the angular resolution may be as large as several areseconds.," Most of the GC X-ray sources in our sample are located away from the aim point, in regions where the angular resolution may be as large as several arcseconds."893 If. therefore. some of the high luminosity sources we detect in bright GCs are actually composites. one or more of the individual components may consist of X-ray sources that have been ejected from the cluster.," If, therefore, some of the high luminosity sources we detect in bright GCs are actually composites, one or more of the individual components may consist of X-ray sources that have been ejected from the cluster."894 It is possible that a subset of such systems could be resolved by pointed observations withChandre's High Resolution Camera., It is possible that a subset of such systems could be resolved by pointed observations with High Resolution Camera.895 The observations we have described introduce a puzzle: the brightest blue GCs do not host X-ray sources. while dimmer blue GCs can house X-ray sources that are among the brightest in the galaxy.," The observations we have described introduce a puzzle: the brightest blue GCs do not host X-ray sources, while dimmer blue GCs can house X-ray sources that are among the brightest in the galaxy."896 We therefore seek to understand how they come to house extremely bright X-ray sources. while the brightest blue GCs do not appear to harbor X-ray sources. even X-ray sources that are ~10 times dimmer.," We therefore seek to understand how they come to house extremely bright X-ray sources, while the brightest blue GCs do not appear to harbor X-ray sources, even X-ray sources that are $\sim 10$ times dimmer."897 The thermal-time-seale scenario described above is likely to work in either young GCs. where the turn-off mass may be larger than 0.8M... or in very massive GCs. in which the probability of a NS orbiting a blue straggler donor is larger.," The thermal-time-scale scenario described above is likely to work in either young GCs, where the turn-off mass may be larger than $0.8\, M_\odot$, or in very massive GCs, in which the probability of a NS orbiting a blue straggler donor is larger."898 Similarly. multiplicity would seem to be most likely in large GCs with a high level of concentration and hence. high central surface brightness.," Similarly, multiplicity would seem to be most likely in large GCs with a high level of concentration and hence, high central surface brightness."899 The problem posed above would seem to require another solution., The problem posed above would seem to require another solution.900 Presumably the donor stars are old. and are less massive than their NS companions.," Presumably the donor stars are old, and are less massive than their NS companions."901 In such cases. the mass transfer rate can be high if the donor ts a giant.," In such cases, the mass transfer rate can be high if the donor is a giant."902 Such systems will exist in a GC only if a binary with total mass ~2.2M... and a separation of ~50—1000R.. can survive.," Such systems will exist in a GC only if a binary with total mass $\sim 2.2\, M_\odot$, and a separation of $\sim 50-1000\, R_\odot$ can survive."903 Some such binaries are considered “soft” in the GC environment. having less binding energy than the ambient kinetic energy of surrounding stars.," Some such binaries are considered “soft” in the GC environment, having less binding energy than the ambient kinetic energy of surrounding stars."904 Soft binaries have a good chance of being disrupted by a combination of distant interactions acting over time., Soft binaries have a good chance of being disrupted by a combination of distant interactions acting over time.905" Others of these binaries may not be formally ""soft"". but. since they are likely to be found near the cluster center. where close interactions are most frequent: they can also be disrupted."," Others of these binaries may not be formally “soft”, but, since they are likely to be found near the cluster center, where close interactions are most frequent; they can also be disrupted."906 Both mechanisms are least efficient at destroying the binaries with large separations in less massive and less highly concentrated GCs., Both mechanisms are least efficient at destroying the binaries with large separations in less massive and less highly concentrated GCs.907 It may therefore be that the most X-ray luminous X-ray sources in the less bright blue GCs are accreting NSs with giant donors., It may therefore be that the most X-ray luminous X-ray sources in the less bright blue GCs are accreting NSs with giant donors.908 It would be difficult to obtain direct evidence that a given GC X-ray source. even in M31. was an example of à NS accreting mass from a giant donor.," It would be difficult to obtain direct evidence that a given GC X-ray source, even in M31, was an example of a NS accreting mass from a giant donor."909 A useful form of indirect evidence could. however. come from GCs in our own Galaxy.," A useful form of indirect evidence could, however, come from GCs in our own Galaxy."910 Specifically. we could attempt to establish the existence. in some clusters with low present-day interaction probabilities. of relatively luminous (L>107 erg s! ) CVs with giant donors.," Specifically, we could attempt to establish the existence, in some clusters with low present-day interaction probabilities, of relatively luminous $L > 10^{32}$ erg $^{-1}$ ) CVs with giant donors."911 Galactic GCs are near enough that the optical properties of the X-ray binaries can be studied withHST. allowing general features of the model to be checked.," Galactic GCs are near enough that the optical properties of the X-ray binaries can be studied with, allowing general features of the model to be checked."912 WDs accreting from giant donors. which cannot presently be studied in external galaxies. may be more common analogs of the accreting NSs with giant donors.," WDs accreting from giant donors, which cannot presently be studied in external galaxies, may be more common analogs of the accreting NSs with giant donors."913 As such. they could provide insight into the environments that nurture such systems.," As such, they could provide insight into the environments that nurture such systems."914 M104 may be one of the richest theatres in which to study X-ray emission from galaxies., M104 may be one of the richest theatres in which to study X-ray emission from galaxies.915 In this paper we have focused on the emission from non-nuclear point sources. since the diffuse emission (see Figure 2) which motivated this observation is the subject of a separate paper (Forman et al.," In this paper we have focused on the emission from non-nuclear point sources, since the diffuse emission (see Figure 2) which motivated this observation is the subject of a separate paper (Forman et al."916 2003: see also Delain et 2001). as is the nuclear source (Pellegrini et 2003).," 2003; see also Delain et 2001), as is the nuclear source (Pellegrini et 2003)."917 The primary science results we have extracted so far from the 18.5 Ksec observation revolve around SSSs and the connection between GCs and LMXBs., The primary science results we have extracted so far from the $18.5$ ksec observation revolve around SSSs and the connection between GCs and LMXBs.918 We find SSSs inhabiting the bulge. disk and halo of M104.," We find SSSs inhabiting the bulge, disk and halo of M104."919 At least 4 of the halo sources are very soft. likely as soft as the SSSs known in the Magellanic Clouds.," At least $4$ of the halo sources are very soft, likely as soft as the SSSs known in the Magellanic Clouds."920 These are good candidates for nuclear-burning accreting WDs., These are good candidates for nuclear-burning accreting WDs.921 The discovery of these sources several kiloparsecs from the disk. indicates that they are old systems. as should be expected for a subset of nuclear-burning WDs.," The discovery of these sources several kiloparsecs from the disk, indicates that they are old systems, as should be expected for a subset of nuclear-burning WDs."922 In the halo as well as in the disk and bulge. we also find SSSs from which we have collected à small number of photons with energies above 1.1 keV. In some cases. photon statistics or the incomplete subtraction of background due to diffuse emission could be responsible for the harder photons.," In the halo as well as in the disk and bulge, we also find SSSs from which we have collected a small number of photons with energies above $1.1$ keV. In some cases, photon statistics or the incomplete subtraction of background due to diffuse emission could be responsible for the harder photons."923 In most cases. however. the photons are likely to emanate from the system itself.," In most cases, however, the photons are likely to emanate from the system itself."924 Some SSSs exhibiting this harder emission may be supernova remnants., Some SSSs exhibiting this harder emission may be supernova remnants.925 Since. however. similar sources have been discovered in other galaxies. including M31. MIOI. NGC 4472. MSI. and M$3. and since we have discovered that many such SSSs are variable on scales of months to years. it is likely that a large fraction are X-ray binaries.," Since, however, similar sources have been discovered in other galaxies, including M31, M101, NGC 4472, M51, and M83, and since we have discovered that many such SSSs are variable on scales of months to years, it is likely that a large fraction are X-ray binaries."926 The harder photons from some SSS binaries may support a model in which the accretor ts an intermediate-mass BH., The harder photons from some SSS binaries may support a model in which the accretor is an intermediate-mass BH.927 We find a significant overdensity of SSSs within 1 kpe of the nucleus: future work will assess, We find a significant overdensity of SSSs within $1$ kpc of the nucleus; future work will assess928component can significantly improve the results.,component can significantly improve the results.929 Even if these models require one more free parameter than the CUTOFFPL with a variable £7... the results we obtained would then be in agreement with those found in the case of TThe spectrum of this source could. indeed. not be reproduced using a simple CUTOFFPL model (see below).," Even if these models require one more free parameter than the CUTOFFPL with a variable $E_{\rm cut}$, the results we obtained would then be in agreement with those found in the case of The spectrum of this source could, indeed, not be reproduced using a simple CUTOFFPL model (see below)."930 A soft spectral component below 2 keV might be expected in the spectra of the SEXT sources. as this component ts very common in binaries hosting a NS accreting mass from a massive companion.," A soft spectral component below 2 keV might be expected in the spectra of the SFXT sources, as this component is very common in binaries hosting a NS accreting mass from a massive companion."931 Hickoxetal.(2004) showed that the detectability of this component depends mainly on the amount of absorption in the direction of the sources., \citet{hickox04} showed that the detectability of this component depends mainly on the amount of absorption in the direction of the sources.932 According to their results. in the most luminous objects (x 107 erg/s) the soft component is produced by reprocessing hard X-rays from the NS by some optically thick material (e.g.. an accretion disk). whereas for sources with lower luminosities (<10°° erg/s) the most likely origin of the soft component is the emission from either a phototonized or collisionally heated diffuse gas in the binary system or from the NS surface.," According to their results, in the most luminous objects $L_{\rm X}$$\simeq$ $^{38}$ erg/s) the soft component is produced by reprocessing hard X-rays from the NS by some optically thick material (e.g., an accretion disk), whereas for sources with lower luminosities $<$ $^{36}$ erg/s) the most likely origin of the soft component is the emission from either a photoionized or collisionally heated diffuse gas in the binary system or from the NS surface."933 In the case of wwe found that BB emission from a relatively cold (—0.1 keV) and large (~ 100 km equivalent radius) region or from a much hotter (Z1 keV) and less extended («100 m) spot. or. alternatively. emission from an optically thin. gas (MKL) provided equally. good fits to the data.," In the case of we found that BB emission from a relatively cold $\sim$ 0.1 keV) and large $\sim$ 100 km equivalent radius) region or from a much hotter $\gtrsim$ 1 keV) and less extended $<$ 100 m) spot, or, alternatively, emission from an optically thin gas (MKL) provided equally good fits to the data."934 While the discussion above and the rapid variability observed in. the SEXT would argue against the presence of an accretion disk in these sources (seealso.Bozzoetal..2008).. we also consider that emission from a small and hot spot on the NS surface is inconsistent. with models of accretion onto magnetic NS.," While the discussion above and the rapid variability observed in the SFXT would argue against the presence of an accretion disk in these sources \citep[see also,][]{bozzo08}, we also consider that emission from a small and hot spot on the NS surface is inconsistent with models of accretion onto magnetic NS."935 The aceretion flow in SEXT is. indeed. thought to be quasi-spherical. and expected to penetrate the NS magnetosphere mainly by means of the Rayleigh-Taylor and the Kelvin-Helmholtz instability (seee.g..Bozzoetal..2008.andref-erences therein)..," The accretion flow in SFXT is, indeed, thought to be quasi-spherical, and expected to penetrate the NS magnetosphere mainly by means of the Rayleigh-Taylor and the Kelvin-Helmholtz instability \citep[see e.g.,][and references therein]{bozzo08}."936 In these circumstances. the size of the hot- over which the accretion takes place is expected to be inversely proportional to the X-ray luminosity and might cover a substantial fraction of the NS surface for Lx X107 erg/s (seeWhiteetal..1983.andreferences therein)..," In these circumstances, the size of the hot-spot over which the accretion takes place is expected to be inversely proportional to the X-ray luminosity and might cover a substantial fraction of the NS surface for $L_{\rm X}$$\lesssim$ $^{35}$ erg/s \citep[see][and references therein]{white83}."937 Taking these results into account. emission from an optically thin and diffuse gas around the NS seems to be a more reasonable explanation of the soft spectral component of According to this interpretation. the emitting region can be estimated from the normalization of the MKL component (see Table 2)) using the relation (seee.g. therein):: Rau VINKL10Llp Ποπ... em. where NAIL is the normalization of the MKL component. D is the source distance. npo πα. e is the binary separation. and a4320/1017 em.," Taking these results into account, emission from an optically thin and diffuse gas around the NS seems to be a more reasonable explanation of the soft spectral component of According to this interpretation, the emitting region can be estimated from the normalization of the MKL component (see Table \ref{tab:xtefit}) ) using the relation \citep[see e.g.,][and references therein]{zurita09b}: $R_{\rm em}$ $\sqrt[3]{3 N_{\rm MKL}/10^{-14} (D/n_{\rm H})^2}$$\sim$ $^{13}$$a_{13}^{2/3}$ cm, where $N_{\rm MKL}$ is the normalization of the MKL component, D is the source distance, $n_{\rm H}$$\sim$$N_{\rm H}$ $a$, $a$ is the binary separation, and $a_{13}$ $a$ $^{13}$ cm."938" Therefore. the radius 72,4, of the emitting region is compatible with the binary separation for a wide range of values of orbital periods similar to those measured in other SFXTs."," Therefore, the radius $R_{\rm em}$ of the emitting region is compatible with the binary separation for a wide range of values of orbital periods similar to those measured in other SFXTs."939 We note also that the properties of this MKL component would be rather similar to those derived from the spectrum of the SFXT 1845.0-0433 Walter. 2009)., We note also that the properties of this MKL component would be rather similar to those derived from the spectrum of the SFXT J1845.0-0433 \citep{zurita09b}.940 Another possibility that we investigated in Sect., Another possibility that we investigated in Sect.941 4.1 is the applicability of a model comprising a power law component and a partial covering to the X-ray spectrum of WWe concluded that this model can also provide a reasonably good fit to the data., \ref{sec:xteresults} is the applicability of a model comprising a power law component and a partial covering to the X-ray spectrum of We concluded that this model can also provide a reasonably good fit to the data.942 A similar model was proposed to interpret the quiescent sspectrum of the SEXT JJ16207-5129 (Tomsicketal..2009) and might provide support for clumpy wind in these sources., A similar model was proposed to interpret the quiescent spectrum of the SFXT J16207-5129 \citep{tomsick09} and might provide support for clumpy wind in these sources.943 In this case. one would expect part of the radiation from the NS to escape absorption by the clumps local to the source and be affected only by interstellar absorption (see.e.g.Walter&ZuritaHeras.2007).," In this case, one would expect part of the radiation from the NS to escape absorption by the clumps local to the source and be affected only by interstellar absorption \citep[see, e.g.][]{walter07}."944. Finally. we showed that the COMPTT model also provided a reasonable fit to the sspectra of aand should thus be considered a valid alternative to. the other models discussed above.," Finally, we showed that the COMPTT model also provided a reasonable fit to the spectra of and should thus be considered a valid alternative to the other models discussed above."945 We note. however. that neither the partial covering nor the COMPTT model could give an acceptable fit to the spectrum of ((see also below).," We note, however, that neither the partial covering nor the COMPTT model could give an acceptable fit to the spectrum of (see also below)."946 Given the similarities between the two sources. a spectral model that provides acceptable results in. both. cases should probably be favored (e.g. the CUTOFFPL+MKL model).," Given the similarities between the two sources, a spectral model that provides acceptable results in both cases should probably be favored (e.g, the CUTOFFPL+MKL model)."947 The oobservation of ddetected a very similar behavior to that discussed above for IIn the light curve of aa number of relatively small flares were observed to take place sporadically on a timescale of few thousands of seconds and were characterized by an X-ray flux — 10-15 times higher than the underlying fainter persistent emission., The observation of detected a very similar behavior to that discussed above for In the light curve of a number of relatively small flares were observed to take place sporadically on a timescale of few thousands of seconds and were characterized by an X-ray flux $\sim$ 10-15 times higher than the underlying fainter persistent emission.948 The lowest X-ray flux that we measured from this source was —3 «10. 1? ere/cm?/s (0.5-10 keV) and corresponds to a luminosity of « 1077 erg/s (at a distance of 3 kpe). comparable to the value reported by Kennea&Cam-pana (2006).," The lowest X-ray flux that we measured from this source was $\sim$ $\times$ $^{-13}$ $^2$ /s (0.5-10 keV) and corresponds to a luminosity of $\times$ $^{32}$ erg/s (at a distance of 3 kpc), comparable to the value reported by \citet{kennea06}."949. The total dynamic range in the X-ray luminosity of bbetween outburst and quiescence is thus 10! (see also Sect. 1)., The total dynamic range in the X-ray luminosity of between outburst and quiescence is thus $\gtrsim$ $^4$ (see also Sect. \ref{sec:intro}) ).950 As for tthe hardness intensity. diagrams and the rate resolved analysis carried out in Sect., As for the hardness intensity diagrams and the rate resolved analysis carried out in Sect.951 4. showed that the variations in the X-ray flux measured during the oobservation were also accompanied by a significant change in the spectral properties of the source., \ref{sec:results} showed that the variations in the X-ray flux measured during the observation were also accompanied by a significant change in the spectral properties of the source.952 In contrast to the case of tthe rate-resolved spectra of ccould not be fit by using a simple CUTOFFPL model., In contrast to the case of the rate-resolved spectra of could not be fit by using a simple CUTOFFPL model.953 We showed that an acceptable fit to the data could. instead. be obtained by introducing an additional relatively cold (KT ~0.08 keV) and extended (2100 km) BB component. or a MKL model (see Tables 3 and 4)).," We showed that an acceptable fit to the data could, instead, be obtained by introducing an additional relatively cold $kT$$\sim$ 0.08 keV) and extended $\gtrsim$ 100 km) BB component, or a MKL model (see Tables \ref{tab:igrbfit} and \ref{tab:igrbfittotal}) )."954 We note that the parameters measured for the MKL component are rather similar to those derived in the case of, We note that the parameters measured for the MKL component are rather similar to those derived in the case of955 (Mileretal.1998).. (Lamb&Miller2001:Abhruuowiezeal.2003:Titarchul:2002).. (Bhattacharvya," \citep[e.g.][]{latt_prak07,klahn06}. \citep{miller98}. \citep{lamb01,stella99,abramowicz03,titarchuk02}. \citep{bhattacharyya05},"956 but suffer frou distance and/or model uncertainties., but suffer from distance and/or model uncertainties.957" Additionally, Nara absorption lines from tje stellar surface would measure the gravitational redshift. but these lines appear to be extremely rare and difficult to observe. and theouly detection is of modest statistical significance. (Cottametal.2002:IKongct 2007)."," Additionally, X-ray absorption lines from the stellar surface would measure the gravitational redshift, but these lines appear to be extremely rare and difficult to observe, and theonly detection is of modest statistical significance \citep{cottam02,kong07}."958. In this paper we prescut a method for placing ταyper limits on ueutron star radii using relativistic won cussion lines (seeMiller2007.forareview] that originate iithe inner accretion disks around the neutron stars in N-ray πως., In this paper we present a method for placing upper limits on neutron star radii using relativistic iron emission lines \citep[see][for a review]{miller07} that originate in the inner accretion disks around the neutron stars in X-ray binaries.959 Velocity shifts eucoded iuto disk eiission lues directly reflect the radius at which the lie is formed., Velocity shifts encoded into disk emission lines directly reflect the radius at which the line is formed.960 Ii the X-ray spectra of accreting black roles and neutron stars. mou Ix-shell eiuissiou lines are widely observed (Naudraetal.etal.1987:Asai 2000).," In the X-ray spectra of accreting black holes and neutron stars, iron K-shell emission lines are widely observed \citep{nandra97,reynolds97,white85,white86,hirano87,asai00}."961. The extreme redwing of sole iron cluission lines nuplies that they are forme: in the immer accretiou disks and are primarily shape x dynamics including relaivistic Doppler shifts. duc to he high velocities in the disk. axd gravitational redshifts (hence the lines are referred to as being relativistic’).," The extreme red-wing of some iron emission lines implies that they are formed in the inner accretion disks and are primarily shaped by dynamics including relativistic Doppler shifts, due to the high velocities in the disk, and gravitational redshifts (hence the lines are referred to as being `relativistic')."962 Therefore. vou Is. lines cau serve as meisive measures of he inner disk radius (Fabianctal.1989).," Therefore, iron K lines can serve as incisive measures of the inner disk radius \citep{fabian89}."963. These relativistic. asvuuimetric Fe Ik emission lines are ποον aud well-studied in the X-ray spectra of yoth supermassive black holes in Active Galactic Nuclei (Naudractal.1997:Revnolds1997). and stellar-nuass lack holes im X-rav binaries (Milleretal.2001).," These relativistic, asymmetric Fe K emission lines are well-known and well-studied in the X-ray spectra of both supermassive black holes in Active Galactic Nuclei \citep{nandra97,reynolds97} and stellar-mass black holes in X-ray binaries \citep{miller04}."964". Iu he case of neutron stars, ron enudssion lines are weaker. lOWOCVOLY. and until lately prior observations have not clearly revealed a relativistic line profile. aud the lines could be well fit by just a Gaussian (e.g.Whiteetal.etal. 2005)."," In the case of neutron stars, iron emission lines are weaker, however, and until lately prior observations have not clearly revealed a relativistic line profile, and the lines could be well fit by just a Gaussian \citep[e.g.,][]{white86,asai00,oosterbroek01,disalvo05}."965. Ouly ταν receutlv have observations of Serpens N-l with shown a relativistic. asvuuuetric line ina neutron star low-mass N-ray binary (LAINB) for the first time (Bhattacharvya&Strolunaver 2001).. confirming the Πιο disk nature of the linesiu," Only very recently have observations of Serpens X-1 with shown a relativistic, asymmetric line in a neutron star low-mass X-ray binary (LMXB) for the first time \citep{bhattacharyya07}, , confirming the inner disk nature of the linesin"966or the particles closest distance to the center of Jupiter became smaller than Jupiter's radius.,or the particle's closest distance to the center of Jupiter became smaller than Jupiter's radius.967" Hlere. ος and λος represent the semimajor axis ancl mass of a Galilean satellite. and AM, is the mass of Jupiter."," Here, $a_{GS}$ and $M_{GS}$ represent the semimajor axis and mass of a Galilean satellite, and $M_J$ is the mass of Jupiter."968 Figure 3 shows the eraphs of the test particles. lifetimes in terms of their initial semümajor axes [or particles in (wo coplanar svstems., Figure 3 shows the graphs of the test particles' lifetimes in terms of their initial semimajor axes for particles in two coplanar systems.969 The eraph al (he top corresponds (o particles initially in circular orbits. aud the one at the bottom shows the liletimes of particles with initial eccentricities of 0.2.," The graph at the top corresponds to particles initially in circular orbits, and the one at the bottom shows the lifetimes of particles with initial eccentricities of 0.2."970 The positions and lifetimes ol the regular satellites of Jupiter and the orbit of Themisto are also shown., The positions and lifetimes of the regular satellites of Jupiter and the orbit of Themisto are also shown.971 As shown bv the upper graph. test particles in circular orbils are mostly stable (for the duration of integrations) except lor a few (hat are close to Callisto.," As shown by the upper graph, test particles in circular orbits are mostly stable (for the duration of integrations) except for a few that are close to Callisto."972 The region of stabilitv. however. becomes smaller (instability progresses (ο. larger distances) in simulations in which the initial eccentricities of test particles are larger.," The region of stability, however, becomes smaller (instability progresses toward larger distances) in simulations in which the initial eccentricities of test particles are larger."973" This can be seen more clearly in figure 4 where from the graplis of figure 3. only the regions between 30/2, and 80/2; are shown."," This can be seen more clearly in figure 4 where from the graphs of figure 3, only the regions between $30{R_J}$ and $80{R_J}$ are shown."974 The islands of instability. corresponding to mean-motion resonances will Callisto (indicated by the subscript C) and Ganvanede (indicated by the subscript G) are also shown.," The islands of instability, corresponding to mean-motion resonances with Callisto (indicated by the subscript C) and Ganymede (indicated by the subscript G) are also shown."975 The migration of unstable regions to larger distances in svstems where test. particles were initially in eccentric orbits was observed in all our simulations., The migration of unstable regions to larger distances in systems where test particles were initially in eccentric orbits was observed in all our simulations.976 Figure 5 shows another exanple of such a system., Figure 5 shows another example of such a system.977 In Chis figure. the lifetimes of test particles with initial eccentricities ol 0.4 and initial inclinations of 20° are shown.," In this figure, the lifetimes of test particles with initial eccentricities of 0.4 and initial inclinations of $^\circ$ are shown."978 The unstable region extends (o distances bevond their corresponding regions in figures 3 and 4., The unstable region extends to distances beyond their corresponding regions in figures 3 and 4.979 We also simulated the dynamics of test particles having orbital inclinations larger than 90° (retrograde orbits)., We also simulated the dynamics of test particles having orbital inclinations larger than $^\circ$ (retrograde orbits).980 As shown by figure 2. the number of irregular satellites is larger al (hese angles implving (hal retrograde orbits have longer lifetimes (HaaniltonanclIxrivov1997:ToumaandWisdom1998:Nesvornyetal. 2003).," As shown by figure 2, the number of irregular satellites is larger at these angles implying that retrograde orbits have longer lifetimes \citep{Hamilton97,Touma98,Nesvorny03}."981. Our simulations also show that retrograde orbits are more stable (han their corresponding prograde ones., Our simulations also show that retrograde orbits are more stable than their corresponding prograde ones.982 Figure 6 shows this lor (wo sels of test particles., Figure 6 shows this for two sets of test particles.983" The particles in black correspond to a svstem in which ey=0.4 and 7,=40°.", The particles in black correspond to a system in which ${e_p}=0.4$ and ${i_p}=40^\circ$.984" The particles in red. correspond to a system with similar orbital eccentricily. but with /,=140."," The particles in red correspond to a system with similar orbital eccentricity, but with ${i_p}=140^\circ$."985 As expected. the particles on retrograde orbits are more stable and maintain their orbits lor longer times.," As expected, the particles on retrograde orbits are more stable and maintain their orbits for longer times."986 The fact that the reeion of instabilitv of test particles. having a given senimajor axis. expands by increasing (he initial values of their orbital eccentriciGes can be attributed to the interactions of (hese particles wilh Jupiters regular satellites.," The fact that the region of instability of test particles, having a given semimajor axis, expands by increasing the initial values of their orbital eccentricities can be attributed to the interactions of these particles with Jupiter's regular satellites."987 Given that the orbits of Jovian regulars are almost circular. an eccentric orbit for a test particle implies a smaller periastron distance for this object. ancl consequently a closer approach to the system's regular moons.," Given that the orbits of Jovian regulars are almost circular, an eccentric orbit for a test particle implies a smaller periastron distance for this object, and consequently a closer approach to the system's regular moons."988 Instability occurs when the perturbative effects of regular satellites disturb the motion of, Instability occurs when the perturbative effects of regular satellites disturb the motion of989As it was noted above. the data of Alcaino’s observations poorly represeut the brightest part of the GB on the Ves.(BVj color - maenitucde diagram.,"As it was noted above, the data of Alcaino's observations poorly represent the brightest part of the GB on the $V vs. (B-V)$ color - magnitude diagram."990 It also refers to the Ves.(BRB) or Ves.(PR) diagrais.," It also refers to the $V vs. (B-R)$ or $V991vs.(V-R)$ diagrams."992 Therefore. the construction of these diagrams and determination for the GB parameters ou them should be carried out as follows.," Therefore, the construction of these diagrams and determination for the GB parameters on them should be carried out as follows."993 From the data obtained by Alcaino et al;," From the data obtained by Alcaino et al.,"994 we derive relations between (2BV) aud (B. R). aud also between (B.V) and (WR) colorsfor cach cluster.," we derive relations between $(B-V)$ and $(B-R)$ , and also between $(B- V)$ and $(V-R)$ colorsfor each cluster."995 The correlation of (2BV) and (P.R) colors is high (>99%). substantially exceeding that in the case of (BW) and (WR) colors. which justifies the selection of the Ves.(BR) diagram for our study.," The correlation of $(B-V)$ and $(B-R)$ colors is high $(>99\%)$, substantially exceeding that in the case of $(B-V)$ and $(V-R)$ colors, which justifies the selection of the $V vs. (B- R)$ diagram for our study."996 With the use of these relations. we transform the Wes.(2BWW) diagrams described in the previous paragraph iutoVWes.(B Ry.," With the use of these relations, we transform the $V997vs.(B-V)$ diagrams described in the previous paragraph into$V998vs.(B-R)$ ."999 When these diagrams are compared with, When these diagrams are compared with1000The CR spectra. further used to calculate the ionization rate. cau be deduced from the measurecl eamuna-ray [lux at Earth: Asstuning that the photons arise from hadronic interactions between the SNR aud the MC. a primary CR spectrum cau be derived.,"The CR spectra, further used to calculate the ionization rate, can be deduced from the measured gamma-ray flux at Earth: Assuming that the photons arise from hadronic interactions between the SNR and the MC, a primary CR spectrum can be derived."1001" We assume a smiootlied broken power-law lor the priuiaries: where αρ is the proton spectrum normalizationaud pp, the location of the spectral break.", We assume a smoothed broken power-law for the primaries: where $a_p$ is the proton spectrum normalizationand $p_{br}$ the location of the spectral break.1002 For all sources. this shape of the primary spectrum reproduces tle shape of the observed photou spectruu.," For all sources, this shape of the primary spectrum reproduces the shape of the observed photon spectrum."1003 Qur results generally follow the original calculations by Fermi. (Abdoetal.2009.2010a.b.c.d). with slight moclilicatious of the smoothing of the power-law. ouly giving nuinor changes to the result.," Our results generally follow the original calculations by Fermi, \citep{abdo_w51c,abdo_w28,abdo_w44,abdo_w49b,adbo_ic443}1004 with slight modifications of the smoothing of the power-law, only giving minor changes to the result."1005 We use the parametrization of the CR interaction model as preseutec in Ixamaeetal.(2006) to derive the primary spectrum from the gamaua-ray observation., We use the parametrization of the CR interaction model as presented in \cite{kamae2006} to derive the primary spectrum from the gamma-ray observation.1006" We use a fixed value of the hydrogeu density. vy,=100 eim7 and determine the CR energy IW), needed to provide the observed flux."," We use a fixed value of the hydrogen density, $n_H=100$ $^{-3}$ and determine the CR energy $W_p$ needed to provide the observed gamma-ray flux."1007 The volume of the emitting region V. enters the normalization of the proton spectrum., The volume of the emitting region $V$ enters the normalization of the proton spectrum.1008 It can be derived (rom gamuua-observatious and is taken from the Fermi publicatious., It can be derived from gamma-observations and is taken from the Fermi publications.1009 The spectral parameters s. As aud ppp. the total energv budget of interacting protous. Hj and the volume V are listed in Table 1..," The spectral parameters $s$, $\Delta s$ and $p_{br}$, the total energy budget of interacting protons, $W_p$ and the volume $V$ are listed in Table \ref{params:tab}."1010 To calculate the ionization rate. we use the parametrization [rom Padovanietal. (2009).. including ionization by CR electrous. CR protous aud electrou capture by CR protons.," To calculate the ionization rate, we use the parametrization from \cite{padovani2009}, , including ionization by CR electrons, CR protons and electron capture by CR protons."1011 Only the, Only the1012An important observational quantity connected to interstellar density and magnetic field fluctuations is the Faraday effect.,An important observational quantity connected to interstellar density and magnetic field fluctuations is the Faraday effect.1013" In the presence of magnetic fields and free electrons, bifringence of circularly polarized orthogonal modes occurs, giving these modes two different propagation velocities."," In the presence of magnetic fields and free electrons, bifringence of circularly polarized orthogonal modes occurs, giving these modes two different propagation velocities."1014" In the case of pure polarized background emission propagation through a magnetoionized medium, the linearly polarized radiation will emerge with its polarization position angle rotated by the amount given in Equation 1.."," In the case of pure polarized background emission propagation through a magnetoionized medium, the linearly polarized radiation will emerge with its polarization position angle rotated by the amount given in Equation \ref{RM}."1015" Thus the relation between the observed position angle, emitted position angle, and the rotation measure is: which has units of radians."," Thus the relation between the observed position angle, emitted position angle, and the rotation measure is: which has units of radians."1016 Here A is the wavelength in meters., Here $\lambda$ is the wavelength in meters.1017" Observational determination of Faraday rotation comes from measurements of the linear polarization vector P= (Q,U) (which depends on Stokes U and Q as |P|=/Q?+ U?) as a function of A?."," Observational determination of Faraday rotation comes from measurements of the linear polarization vector $\textbf{P} \equiv$ (Q,U) (which depends on Stokes U and Q as $|P|=\sqrt{Q^2+U^2}$ ) as a function of $\lambda^2$."1018" To avoid confusion between vector and scalar P, we use bold notation to denote the vector quantity of the linear polarization map."," To avoid confusion between vector and scalar P, we use bold notation to denote the vector quantity of the linear polarization map."1019" We define the gradient of the polarization vector as: We note that in the case of vector P we have: From this equation, one can derive a relationship between |VP| and |VRM| for Faraday-thin polarized emission as: Equation 5 only holds for data in which the entire signal is measured (i.e. single-dish data included) and for which the background is uniform."," We define the gradient of the polarization vector as: We note that in the case of vector $\textbf{P}$ we have: From this equation, one can derive a relationship between $|\nabla \textbf{P}|$ and $|\nabla RM|$ for Faraday-thin polarized emission as: Equation 5 only holds for data in which the entire signal is measured (i.e. single-dish data included) and for which the background is uniform."1020" When |P|=1 and \=1 (which are the assumptions we use for the simulations), one finds a trivial relation between |VP| and |VRM| as |\VRM|=|VP|/2i."," When $|\textbf{P}|=1$ and $\lambda =1$ (which are the assumptions we use for the simulations), one finds a trivial relation between $|\nabla \textbf{P}|$ and $|\nabla RM|$ as $|\nabla RM|= |\nabla \textbf{P}|/2i$."1021" However, this relation can only be used to calculate |VP| or |VRM| in the simulations, since the assumption of |P|=1 is almost always to simplistic for the observations because the data are missing single-dish information and/or the background |VP| is not zero."," However, this relation can only be used to calculate $|\nabla \textbf{P}|$ or $|\nabla RM|$ in the simulations, since the assumption of $|\textbf{P}|=1$ is almost always to simplistic for the observations because the data are missing single-dish information and/or the background $|\nabla \textbf{P}|$ is not zero."1022" We calculate maps of RM, P and their gradients from density and LOS magnetic field that is perpendicular and parallel to the mean magnetic field in the simulations."," We calculate maps of RM, $\textbf{P}$ and their gradients from density and LOS magnetic field that is perpendicular and parallel to the mean magnetic field in the simulations."1023" We calculate the RM as per Equation 1 at every point and then take its spatial gradient, that is, we compute the gradient vector at every pixel of the image using neighbor pixels."," We calculate the RM as per Equation \ref{RM} at every point and then take its spatial gradient, that is, we compute the gradient vector at every pixel of the image using neighbor pixels."1024 We can calculate [P| by calculating the stokes vectors as Q—cos(20) U—sin(20)., We can calculate $|\textbf{P}|$ by calculating the stokes vectors as $Q=\cos(2\theta)$ $U=\sin(2\theta)$.1025" These expressions come from applying Equation 2 with assumed values for λ and 0o. We show a subsonic and supersonic case of density (n), Vn, LOS magnetic field (LOS B), VB, RM and |VP| in Figure 2 and Figure 3,, respectively."," These expressions come from applying Equation 2 with assumed values for $\lambda$ and $\theta_0$ We show a subsonic and supersonic case of density $n$ ), $\nabla n$, LOS magnetic field (LOS B), $\nabla B$, RM and $|\nabla P|$ in Figure \ref{fig:RM1} and Figure \ref{fig:RM2}, respectively."1026 A comparison of the SGPS test data and a subsonic case is given in Figure 1.., A comparison of the SGPS test data and a subsonic case is given in Figure \ref{fig:RM}.1027" Inspection of maps of |VP| reveals that filaments are created in both cases and that there is some correlation between gradients of column density, magnetic field, and |VP|."," Inspection of maps of $|\nabla \textbf{P}|$ reveals that filaments are created in both cases and that there is some correlation between gradients of column density, magnetic field, and $|\nabla \textbf{P}|$."1028 We will discuss these further in Section 3.., We will discuss these further in Section \ref{origin}.1029" We use a subsection of radio continuum images of an 18-square-degree patch of the Galactic plane, observed with the Australia Telescope Compact Array (ATCA, see McClure-Griffiths et al."," We use a subsection of radio continuum images of an 18-square-degree patch of the Galactic plane, observed with the Australia Telescope Compact Array (ATCA, see McClure-Griffiths et al."1030 2001 and Gaensler et al., 2001 and Gaensler et al.1031 2001 for more details)., 2001 for more details).1032" We examine the 1.4 GHz frequency data averaged over adjoining frequency channels with simultaneously recorded StokesJ,, StokesQ,, and Stokes as part of the SGPS test region (Gaensler et al."," We examine the 1.4 GHz frequency data averaged over adjoining frequency channels with simultaneously recorded Stokes, Stokes, and Stokes as part of the SGPS test region (Gaensler et al."1033 , 2001).1034"This field consists of 190 mosaicked pointings of the 2001).Australia Telescope Compact Array (ATCA) and covers the range 325.5<1332.5,-0.5«b 3.5."," This field consists of 190 mosaicked pointings of the Australia Telescope Compact Array (ATCA) and covers the range $325.5 < l < 332.5, -0.5 < b < 3.5$ ."1035" Complicated extended structure is seen in linear polarization throughout the test region, almost all of which has no correlation with total intensity."," Complicated extended structure is seen in linear polarization throughout the test region, almost all of which has no correlation with total intensity."1036 We select a 512x pixel subregion from this data to match the resolution of the simulations used in our study and display it in Figure 1 in the top row., We select a $512\times512$ pixel subregion from this data to match the resolution of the simulations used in our study and display it in Figure \ref{fig:RM} in the top row.1037" The SGPS region we select begins at coordinate 1=332.3373, b= -0.3138 and is not overly contaminated by bad pixels and contains significant emission."," The SGPS region we select begins at coordinate l=332.3373, b= -0.3138 and is not overly contaminated by bad pixels and contains significant emission."1038 We generate a database of 3D numerical simulations of isothermal compressible (MHD) turbulence by using the MHD code of Cho Lazarian (2003) and varying the, We generate a database of 3D numerical simulations of isothermal compressible (MHD) turbulence by using the MHD code of Cho Lazarian (2003) and varying the1039Close sun-grazing comets are discovered almost daily by white light coronographs (e.g. SOHO LASCO).,Close sun-grazing comets are discovered almost daily by white light coronographs (e.g. SoHO LASCO).1040 Most are small and fully sublimated by isolation at a few solar radii Λο while almost none have re-emerged (Marsden 2005)., Most are small and fully sublimated by insolation at a few solar radii $R_\odot$ while almost none have re-emerged (Marsden 2005).1041 The majority have perihelion distances g well above Ra but some have το <Rx» (e.g. Marsden 2005 and http://www.minorplanetcenter.org/mpec/RecentMPECSs.html.), The majority have perihelion distances $q$ well above $R_\odot$ but some have $q\approx$ or $\prec R_\odot$ (e.g. Marsden 2005 and http://www.minorplanetcenter.org/mpec/RecentMPECs.html.)1042 The death of a comet at ry~R« has been seen directly only very recently (Schrijver et al 2011) using the SDO AIA XUV instrument., The death of a comet at $r\sim R_\odot$ has been seen directly only very recently (Schrijver et al 2011) using the SDO AIA XUV instrument.1043 This recorded sublimative destruction of Comet C/2011 N3 as it crossed the solar disk very near perihelion g=L1I39R.., This recorded sublimative destruction of Comet C/2011 N3 as it crossed the solar disk very near perihelion $q=1.139R_\odot$.1044" The next challenge i studying the demise of close sun-grazers will be to cateh one of the even rarer cases of chromospheric impact (g<1.0IRs. M,>10!! σι."," The next challenge in studying the demise of close sun-grazers will be to catch one of the even rarer cases of chromospheric impact $q<1.01R_\odot$, $M_o>10^{11}$ g)."1045 As anticipated by Weissman (1983) and shown below. these undergo explosive destruction in the dense chromosphere.," As anticipated by Weissman (1983) and shown below, these undergo explosive destruction in the dense chromosphere."1046 Understanding the destruction processes. and their radiation signatures. are essential steps in searching for and modeling these.," Understanding the destruction processes, and their radiation signatures, are essential steps in searching for and modeling these."1047 The processes leading to sublimation of the icy conglomerate mix (Whipple 1950) of cometary nuclei. and in some cases their splitting and fragmentation. were considered by Huebner (1967). Weissman and Kieffer (1981). Weissman (1983). Iselt et al. (," The processes leading to sublimation of the icy conglomerate mix (Whipple 1950) of cometary nuclei, and in some cases their splitting and fragmentation, were considered by Huebner (1967), Weissman and Kieffer (1981), Weissman (1983), Iseli et al. ("10482002). Sekanina (2003) and others.,"2002), Sekanina (2003) and others."1049 These models essentially solve for the insolative sublimation mass loss rates of icy-conglomerate mixes (the dust being carried away in the flow of these evolved components)., These models essentially solve for the insolative sublimation mass loss rates of icy-conglomerate mixes (the dust being carried away in the flow of these evolved components).1050 They variously allow for the complicating factors of rotation. albedo. insulating surface dust layers. radiative cooling. interior thermal conduction. and fragmentation by tidal. thermal and volatile explosion effects.," They variously allow for the complicating factors of rotation, albedo, insulating surface dust layers, radiative cooling, interior thermal conduction, and fragmentation by tidal, thermal and volatile explosion effects."1051 Huebner (1967) and Iselt et al. (, Huebner (1967) and Iseli et al. (10522002). for example. found that. for high sublimation rates near the sun. these effects were secondary. the mass loss being reasonably approximated by a pure sublimation description: mass loss rate = heating rate/latent heat.,"2002), for example, found that, for high sublimation rates near the sun, these effects were secondary, the mass loss being reasonably approximated by a pure sublimation description: mass loss rate = heating rate/latent heat."1053 Using SoHO data. Sekanina (2003) addressed in detail how mass loss rates and fragmentation relate to cometary light curves via atomic line emission (e.g. by sodium) and by dust scattering of sunlight. though emphasising that most of the mass remains in a primary fragment.," Using SoHO data, Sekanina (2003) addressed in detail how mass loss rates and fragmentation relate to cometary light curves via atomic line emission (e.g. by sodium) and by dust scattering of sunlight, though emphasising that most of the mass remains in a primary fragment."1054" None of these studies considered ablation or ram-pressure driven explosion due to solar atmospheric impact (which we show below are negligible till 7<1.018.) though Weissman (1983) had remarkedatmosphere"".", None of these studies considered ablation or ram-pressure driven explosion due to solar atmospheric impact (which we show below are negligible till $r\preceq1.01R_\odot$ ) though Weissman (1983) had remarked.1055 It is well known in the planetary physics community that ablation and explosion are central processes in comet-planetary atmosphere impacts (e.g. Carslon et al., It is well known in the planetary physics community that ablation and explosion are central processes in comet-planetary atmosphere impacts (e.g. Carslon et al.1056 1997) Here we revisit the theory of sun-grazer sublimation then develop the first estimates of the much higher rates of mass loss by ablation for the rarer cases which reach <1.01R.., 1997) Here we revisit the theory of sun-grazer sublimation then develop the first estimates of the much higher rates of mass loss by ablation for the rarer cases which reach $\preceq1.01R_\odot$.1057 This regime (Sections 3.4. 6.1. 6.2) resembles that of comet-planetary impacts though with some differences.," This regime (Sections 3.4, 6.1, 6.2) resembles that of comet-planetary impacts though with some differences."1058 One is that almost all sun-grazing comets belong to secondary comet groups (mostlyKreutz) formed from primordial comets., One is that almost all sun-grazing comets belong to secondary comet groups (mostlyKreutz) formed from primordial comets.1059" A significant number of the latter must have g<R« (Hughes 2001) and large enough mass (M,=10! & - see Section 5. Eqn. (27)))"," A significant number of the latter must have $q<R_\odot$ (Hughes 2001) and large enough mass $M_o\succeq 10^{11}$ g - see Section 5, Eqn. \ref{Mominq}) ))"1060 to survive sublimation down to the intense ablation/explosion regime., to survive sublimation down to the intense ablation/explosion regime.1061 However only a small fraction of group comets (Bresecker et al., However only a small fraction of group comets (Biesecker et al.1062 2003. Knight et al 2010) come that close or are that massive. larger sun-grazers mostly having g=VSR...," 2003, Knight et al 2010) come that close or are that massive, larger sun-grazers mostly having $q\succeq 1.5 R_\odot$."1063 Comet C/2010 E6 (STEREO) discussed by Raftery et al. (, Comet C/2010 E6 (STEREO) discussed by Raftery et al. (10642010) came close. having g=1.02R.. while the destruction of Comet C/2011 N3 seen by SDO (Schrijver et al 2011) on the solar disk was solely by sublimation as it had g= Ι.Ν.,"2010) came close, having $q\approx 1.02 R_\odot$ while the destruction of Comet C/2011 N3 seen by SDO (Schrijver et al 2011) on the solar disk was solely by sublimation as it had $q\approx 1.14R_\odot$ ."1065 Most of the few group comets reaching, Most of the few group comets reaching1066The EBL intensity at the present epoch (2=0) provides an inteeral constraint on the history of clectromaeuctic energv release in the universe since recombination.,The EBL intensity at the present epoch $\left(z=0\right)$ provides an integral constraint on the history of electromagnetic energy release in the universe since recombination.1067 Measureineut of this cumulative output however. cannot address its evolution and thus caunot be related to issues such as the history of star iux clement formation.," Measurement of this cumulative output however, cannot address its evolution and thus cannot be related to issues such as the history of star and element formation."1068 For this reason. several models have been developed to calculate the EBL huuinosity density as a function of redshift. £772). from fundaiueuta astroplivsical principles.," For this reason, several models have been developed to calculate the EBL luminosity density as a function of redshift, $\mathcal{L}\left(\nu,z\right)$, from fundamental astrophysical principles."1069 These models were composcc with varvine degrees of complexity. obscrvationa constraints. and data mputs.," These models were composed with varying degrees of complexity, observational constraints, and data inputs."1070 The caleulation of the high-encrey absorption feature iu the blazar componcut of the ECRD requires a mode of the EBL aud its evolution over cosmüc time., The calculation of the high-energy absorption feature in the blazar component of the EGRB requires a model of the EBL and its evolution over cosmic time.1071" The eroat deeree of uncertainty associated with the EBL models aud their predictions reuders sclecting the ""best? candidate model impossible.", The great degree of uncertainty associated with the EBL models and their predictions renders selecting the “best” candidate model impossible.1072 Therefore. we use several models with widely raneine precictious iu order to bracket the rauge of possible EBL realizations.," Therefore, we use several models with widely ranging predictions in order to bracket the range of possible EBL realizations."1073 Isneiske et al. (, Kneiske et al. (10742001) treat the EBL-nodeling problem. using scparate approaches at the UWoptical aac infrared wavelengths.,2004) treat the EBL-modeling problem using separate approaches at the UV–optical and infrared wavelengths.1075 Iu determining the EBL at various wavelengths. they iake use of a cosmic chemica evolution model at the UWoptical wavelengths aud a backwards evolution model for the infrared. (see Taser Dwek (2001) for a complete review of the differeu types of models).," In determining the EBL at various wavelengths, they make use of a cosmic chemical evolution model at the UV–optical wavelengths and a backwards evolution model for the infrared (see Hauser Dwek (2001) for a complete review of the different types of models)."1076 Additionally. this lwbrid model was paralucterized in terms of the iain observationa nucertainties such as the redshift dependence of the cosmic star formation rate and the fraction of UV radiation released frou star formune regions.," Additionally, this hybrid model was parameterized in terms of the main observational uncertainties such as the redshift dependence of the cosmic star formation rate and the fraction of UV radiation released from star forming regions."1077 Thus. the Ixuciske et al.," Thus, the Kneiske et al."1078 EBL 1nodel cousists of several that allow for the iuclusion of various EBL scenarios that are consistent with the available data., EBL model consists of several that allow for the inclusion of various EBL scenarios that are consistent with the available data.1079 Specifically. the model best interpolates the data with the portant caveat that the assumed complete UW absorption by interstellar gas introduces a sharp cutoff at 0.112. In theS," Specifically, the model best interpolates the data with the important caveat that the assumed complete UV absorption by interstellar gas introduces a sharp cutoff at $0.1\mu {\rm m}$."1080tellar- model. all the UV. radiation produced by the stellar populationsUV escapes to the intergalactic medina after reprocessing by the interstellar gas. with theStellaa- model allowing for a strong UV-field at high redshifts.," In the model, all the UV radiation produced by the stellar populations escapes to the intergalactic medium after reprocessing by the interstellar gas, with the model allowing for a strong UV-field at high redshifts."1081 UVSince the x-ray sources likely respousible for the EGRB at GeV euergies are particularly sensitive to the EBL density at UV wavelengths. for the purposes of this analysis. the audHigh-Stellar- models are used to bracket the possible ranges of attenuation.," Since the $\gamma$ -ray sources likely responsible for the EGRB at GeV energies are particularly sensitive to the EBL density at UV wavelengths, for the purposes of this analysis, the and models are used to bracket the possible ranges of attenuation."1082UV. Primack et al. (, Primack et al. (10832008) lave piouecred the use of senianalvtical models that attempt to reproduce the process of structure formation and evolution through sinulatious.,2008) have pioneered the use of semianalytical models that attempt to reproduce the process of structure formation and evolution through simulations.1084 Recent iterations of this model incorporate liehlv precise knowledge of the local Iuninositv deusity i opticalUV (Gilmore et al., Recent iterations of this model incorporate highly precise knowledge of the local luminosity density at optical–UV (Gilmore et al.1085 2009) aud NIR (Primack et al., 2009) and NIR (Primack et al.1086 2008) waveleneths aud a well-established cosmological model., 2008) wavelengths and a well-established cosmological model.1087 The kev parameters in their approach (those that govern the rate of star formation. supernova feedback. and inetalliitv) have been adjusted to fit the local ealaxvo data.," The key parameters in their approach (those that govern the rate of star formation, supernova feedback, and metallicity) have been adjusted to fit the local galaxy data."1088 With respect to estinates bv Primack. et al., With respect to estimates by Primack et al.1089 from previous vears. this version of the model vields a lower huninosity deusitv at optical waveleneths. thereby resulting iu a reduced EBL density.," from previous years, this version of the model yields a lower luminosity density at optical wavelengths, thereby resulting in a reduced EBL density."1090 Receut TeV observations of ucarby blazars secum to support stich low values (Aharouian et al., Recent TeV observations of nearby blazars seem to support such low values (Aharonian et al.1091 2006)., 2006).1092" Finally, we consider the nost recent EBL model bv Stecker et al. ("," Finally, we consider the most recent EBL model by Stecker et al. ("10932006).,2006).1094 Iu this inodoel. Stecker et al.," In this model, Stecker et al."1095 calculate. the EBL at infrared and opticalUV waveleugths separately., calculate the EBL at infrared and optical–UV wavelengths separately.1096 At infrared waveleneths. they use a backavards evolution model based ou observational shnowledge of: (1) huninositv-depeudeut galaxy SEDs. (2) galaxy huninosity functions. and (3) parameterized Unefious for luninesitv evolution.," At infrared wavelengths, they use a backwards evolution model based on observational knowledge of: (1) luminosity-dependent galaxy SEDs, (2) galaxy luminosity functions, and (3) parameterized functions for luminosity evolution."1097 For opticalUV wavelengths. they consider the redshift evolution of stellar populations with an analytical approximation o the more sophisticated SEDs used iu Salamiou Stecker (1998).," For optical–UV wavelengths, they consider the redshift evolution of stellar populations with an analytical approximation to the more sophisticated SEDs used in Salamon Stecker (1998)."1098" The SEDs adapted frou: Druzual Charlot (1993) reflect stellar population svutlesis models or Galaxy evolution aud the observational fact that star forming galaxies are ""bluer (brighter iu the blue wart of the optical spectrum) atf +>0.7.", The SEDs adapted from Bruzual Charlot (1993) reflect stellar population synthesis models for galaxy evolution and the observational fact that star forming galaxies are “bluer” (brighter in the blue part of the optical spectrum) at $z>0.7$.1099 Notably. he UW spectra for all SEDs are assuued to cut off at the Lyman lint. aud the effects of extinction by dust are not included im the model.," Notably, the UV spectra for all SEDs are assumed to cut off at the Lyman limit, and the effects of extinction by dust are not included in the model."1100 The former is a matter of debate since it is not really known how much UV radiation short of the Lyman limit cau escape from the star-formine regions. while the latter would result inexorably in an overpreciction of the UV photou density and. consequently. the optical depth at ligher redshifts.," The former is a matter of debate since it is not really known how much UV radiation short of the Lyman limit can escape from the star-forming regions, while the latter would result inexorably in an overprediction of the UV photon density and, consequently, the optical depth at higher redshifts."1101 In a siuiluw vein. Franceschini et al. (," In a similar vein, Franceschini et al. ("1102200s) also eniploy. a backwards evolution model based rather detailed observations.,2008) also employ a backwards evolution model based rather detailed observations.1103 However. their deteriuination of the EBL departs significantly from that of Stecker et al.," However, their determination of the EBL departs significantly from that of Stecker et al.,"1104 particularly at the optical aud UV waveleneths to which GeV photous are most seusitive., particularly at the optical and UV wavelengths to which GeV photons are most sensitive.1105 The differences iu these models are likely due to differences in the treatiueut of ealaxy evolution., The differences in these models are likely due to differences in the treatment of galaxy evolution.1106 EBL atteuuatiou is a function of the observed οταν enerev £y aud the redshift : of the oemüttiug source., EBL attenuation is a function of the observed $\gamma$ -ray energy $E_0$ and the redshift $z$ of the emitting source.1107 The atteuuation is ecucrally parameterized by the optical depth r(Eg.τὸν which is defined as the πα: of e-fold reductions of the observed flux. £j. as compared witli the emitted source flux. Fue: at redshift <: The optical depth is caleulated from plivsical principles.," The attenuation is generally parameterized by the optical depth $\tau\left(E_0,z\right)$, which is defined as the number of e-fold reductions of the observed flux, $F_{\mathrm{obs}}$, as compared with the emitted source flux, $F_{\mathrm{emitted}}$, at redshift $z$ The optical depth is calculated from physical principles."1108 Using the cross section for pair production c. aud assunuius isotropic backeround radiation with spectral deusitv »(e) at energv €. the absorption probability of 2-xavs per uuit path is given by where @ is the scattering augle for the >5 collision. egyδηοἱ[Eο is the eucrgv threshold for the reaction. and i» is the electron mass.," Using the cross section for pair production $\sigma$, and assuming isotropic background radiation with spectral density $n\left(\epsilon\right)$ at energy $\epsilon$, the absorption probability of $\gamma$ -rays per unit path is given by where $\theta$ is the scattering angle for the $\gamma-\gamma$ collision, $\epsilon_{\rm th}=2m^{2}c^{4}/\left[E\left(1-\cos\theta\right)\right]$ is the energy threshold for the reaction, and $m$ is the electron mass."1109 Since blazars are the sources Όσιο considered. redshift is a eood choice to nicasure the distance. with the total distance beiug the look-back time (times the speed of light. ο) ," Since blazars are the sources being considered, redshift is a good choice to measure the distance, with the total distance being the look-back time (times the speed of light, $c$ "1110"where ©=AVTint)”. which is an estimate of the(CN, noiseN20 powerNN) in the reference cross-correlation. and the prime indicates that C,» has been approximated.","where $\Psi=\left(\left<N^{}_{r_1}N^{*}_{r_1}\right>1111\left<N^{}_{r_2}N^{*}_{r_2}\right>/\Delta\nu\tau_{\it{int}}\right)^{1/2}$, which is an estimate of the noise power in the reference cross-correlation, and the prime indicates that $C_{\it{pc2}}$ has been approximated."1112" The extra term V? stops the zero-mean fluctuations in P.,,, from going too close to zero. while for large INP, equation (18)) reduces to (16))."," The extra term $\Psi^2$ stops the zero-mean fluctuations in $P^{}_{r_1r_2}$ from going too close to zero, while for large $\it{INR}_r$ equation \ref{MODIFIED SPECTRUM1113MK2}) ) reduces to \ref{CORRECTION SPECTRUM MK2}) )."1114 Since V? introduces a small bias in a similar way to the MK] cancellers. a relation equivalent to (7)). but with a much smaller bias. can be derived: This will have both a noise and a RFI component. as in (7)). however. since the canceller is now biased like a MKla canceller the added receiver noise term will not average away.," Since $\Psi^2$ introduces a small bias in a similar way to the MK1 cancellers, a relation equivalent to \ref{mk1 residual power}) ), but with a much smaller bias, can be derived: This will have both a noise and a RFI component, as in \ref{mk1 residual1115power}) ), however, since the canceller is now biased like a MK1a canceller the added receiver noise term will not average away."1116 Hm.» reduces as INR* multiplied by the number of samples in the time average., $R_{\it{pc2}}^{\prime}$ reduces as $\it{INR}^2$ multiplied by the number of samples in the time average.1117 There is a similar problem for the pre-correlation MK2 adaptive cancellers., There is a similar problem for the pre-correlation MK2 adaptive cancellers.1118 In the lag domain. the division in (12)) becomes a multiplication by the inverse of a matrix with columns containing offset copies of the 54-7» cross-correlation function. (Widrow&Stearns1985).," In the lag domain, the division in \ref{mk2 pre corr weights}) ) becomes a multiplication by the inverse of a matrix with columns containing offset copies of the $r_1$ $r_2$ cross-correlation function \citep{Widrow1985}."1119. Divisions by zero in the frequency domain due to interference-free frequency channels in are manifest in the reference lag natrix as singular values., Divisions by zero in the frequency domain due to interference-free frequency channels in $P^{}_{r_1r_2}$ are manifest in the reference lag matrix as singular values.1120" P,..,..One method of dealing with this is to use singular value decomposition to decompose the matrix into two orthonormal triangular matrices and one diagonal natrix (see 5 2.9 of Pressetal. 1986)).", One method of dealing with this is to use singular value decomposition to decompose the matrix into two orthonormal triangular matrices and one diagonal matrix (see $\S$ 2.9 of \citealt{Press1986}) ).1121 Singular (or near-singular) values can be selected when the relevant diagonal natrix elements are less than a chosen threshold. such as V.," Singular (or near-singular) values can be selected when the relevant diagonal matrix elements are less than a chosen threshold, such as $\Psi$."1122 The singular parts of the matrix contain no information about the correlated signal and are removed from the decomposition natrices., The singular parts of the matrix contain no information about the correlated signal and are removed from the decomposition matrices.1123 The inverse matrix can then be reconstructed from the remaining parts of the three decomposition matrices. and it will not function in the RFl-free parts of the spectrum.," The inverse matrix can then be reconstructed from the remaining parts of the three decomposition matrices, and it will not function in the RFI-free parts of the spectrum."1124 Interference cancelling using a single reference antenna can give excellent results when the reference signal interference-to-noise ratio is large. and there is more gain towards the interfering signal for the reference antenna than for the astronomy antennae.," Interference cancelling using a single reference antenna can give excellent results when the reference signal interference-to-noise ratio is large, and there is more gain towards the interfering signal for the reference antenna than for the astronomy antennae."1125 However. receiver noise in the reference signal means that a fraction of the interference will always remain after cancelling.," However, receiver noise in the reference signal means that a fraction of the interference will always remain after cancelling."1126 A second reference signal can be used to remove the noise bias and give infinite interference attenuation. but a larger amount of reference receiver noise is added during cancelling.," A second reference signal can be used to remove the noise bias and give infinite interference attenuation, but a larger amount of reference receiver noise is added during cancelling."1127 For pre-correlation systems. a dual canceller setup can be used to average the (zero-mean) receiver noise away. à process that comes automatically with post-correlation cancellers.," For pre-correlation systems, a dual canceller setup can be used to average the (zero-mean) receiver noise away, a process that comes automatically with post-correlation cancellers."1128 A breakdown of the main properties for the different mitigation techniques is given in table 2.., A breakdown of the main properties for the different mitigation techniques is given in table \ref{SUMMARY TABLE}.1129 It is important to note that even though the single-reference cancellers leave residual RFI. the residual may be extremely small and well below the RMS noise.," It is important to note that even though the single-reference cancellers leave residual RFI, the residual may be extremely small and well below the RMS noise."1130 This occurs when the interference-to-noise ratios of the reference signals are very large. and the use of two references (1n pre-correlation systems) might just add complexity to the system with little or no benefit.," This occurs when the interference-to-noise ratios of the reference signals are very large, and the use of two references (in pre-correlation systems) might just add complexity to the system with little or no benefit."1131 However. if maximum sensitivity is required. one should be aware that they will eventually reach a non-zero residual signal.," However, if maximum sensitivity is required, one should be aware that they will eventually reach a non-zero residual signal."1132 Using two reference signals to remove the reference receiver. noise bias removes the inherent stability of the algorithms in. situations where some or all of the frequency channels are interference-free., Using two reference signals to remove the reference receiver noise bias removes the inherent stability of the algorithms in situations where some or all of the frequency channels are interference-free.1133 Although there are applications in which the passband will always be entirely filled with RFI (such as observations in the GPS LI and L2 bands). many interfering signals will only take up a part of the band.," Although there are applications in which the passband will always be entirely filled with RFI (such as observations in the GPS L1 and L2 bands), many interfering signals will only take up a part of the band."1134 In these cases the algorithms need a mechanism to turn themselves off in the vacant frequency channels., In these cases the algorithms need a mechanism to turn themselves off in the vacant frequency channels.1135The fraction of galaxies m substructure increases with radius in. both ensemble clusters.,The fraction of galaxies in substructure increases with radius in both ensemble clusters.1136 Most of the galaxies belonging to substructures (especially those detected with a global value of 6.) are located at r»rsoo., Most of the galaxies belonging to substructures (especially those detected with a global value of $\delta_{c}$ ) are located at $>$ $_{200}$.1137 The fraction of galaxies detected in substructure is higher in EC? than in EC]., The fraction of galaxies detected in substructure is higher in EC2 than in EC1.1138 Nevertheless. we have seen no trend between the fraction of galaxies in. substructure and their absolute magnitude.," Nevertheless, we have seen no trend between the fraction of galaxies in substructure and their absolute magnitude."1139" Galaxies brighter than M,=—22 are preferentially located outside substructures.", Galaxies brighter than $M_{r}=-22$ are preferentially located outside substructures.1140 Those located 1n substructures (only a few percent) are located 1n the outer regions of the clusters. indicating. that they. could be located in galaxy groups or clusters in the process of merging with the cluster itself.," Those located in substructures (only a few percent) are located in the outer regions of the clusters, indicating that they could be located in galaxy groups or clusters in the process of merging with the cluster itself."1141 We have also investigated the dynamies of the galaxies selected in substructures., We have also investigated the dynamics of the galaxies selected in substructures.1142 Independent of the method usec for selecting galaxies in substructures. they represent a mixed population of backsplash and infalling galaxies.," Independent of the method used for selecting galaxies in substructures, they represent a mixed population of backsplash and infalling galaxies."1143" The substructure galaxies selected by a global value of 6, turnec to be dominated by an infalling population of galaxies.", The substructure galaxies selected by a global value of $\delta_{c}$ turned to be dominated by an infalling population of galaxies.1144" In contrast. those galaxies selected η substructures using individial values of 6, for each cluster turned to be dominated by back-splash galaxies."," In contrast, those galaxies selected in substructures using individial values of $\delta_{c}$ for each cluster turned to be dominated by back-splash galaxies."1145 Assuming that all backsplash galaxies located in substructures are red ones. the fraction of blue galaxies of the infalling population is similar to that observed for isolated objects.," Assuming that all backsplash galaxies located in substructures are red ones, the fraction of blue galaxies of the infalling population is similar to that observed for isolated objects."1146 This indicates that we can not ruled out the hypothesis that the infall population of galaxies located i substructures would be genuine field ones., This indicates that we can not ruled out the hypothesis that the infall population of galaxies located in substructures would be genuine field ones.1147we &o back in cosmic time and that the estimate of number counts based on the HI mass function measured at z=0 is likely to be a lower limit of the number counts found bv a particular survey.,we go back in cosmic time and that the estimate of number counts based on the HI mass function measured at $z=0$ is likely to be a lower limit of the number counts found by a particular survey.1148 We will make a simple assumption that the relevant galaxies are formed by the epochs probed by FA and that HIE is being usec as the fuel for star formation. and hence is reducing in each of the galaxies as a function of time.," We will make a simple assumption that the relevant galaxies are formed by the epochs probed by FAST, and that HI is being used as the fuel for star formation, and hence is reducing in each of the galaxies as a function of time."1149 We note that this passive star formation assumption ignores the possibility of mergers., We note that this passive star formation assumption ignores the possibility of mergers.1150 We further note that this is only à toy model anc is of course [ar too simplistic to be considered. a realistic portraval of the intricate physical processes acting on the LIE mass function., We further note that this is only a toy model and is of course far too simplistic to be considered a realistic portrayal of the intricate physical processes acting on the HI mass function.1151 Lt is designed solely as an illustrative example., It is designed solely as an illustrative example.1152" Η the LL mass function. is well-represented by. à Schechter function then the total density of HE is given by lÜ we now make the assumption that the faint-cnc slope a does not evolve with redshift. and that 8*=0(z)ὃν we can model the ellective change in the hydrogen mass due to passive star formation as a net shift of the break Aj,=Adj(2) with redshift."," If the HI mass function is well-represented by a Schechter function then the total density of HI is given by If we now make the assumption that the faint-end slope $\alpha$ does not evolve with redshift, and that $\theta^*\neq\theta^*(z)$, we can model the effective change in the hydrogen mass due to passive star formation as a net shift of the break $M_{\rm HI}^*=M_{\rm HI}^*(z)$ with redshift."1153 We assume that a simple form of this evolution where 3 is some constant of proportionality that determines the rate at which the hydrogen mass is consumed. by star formation., We assume that a simple form of this evolution where $\beta$ is some constant of proportionality that determines the rate at which the hydrogen mass is consumed by star formation.1154 This should be valid. at low redshifts., This should be valid at low redshifts.1155 In. the subsequent discussion. we will assume that 0x3 in this moclel., In the subsequent discussion we will assume that $0\le\beta\le 3$ in this model.1156 One can determine values compatible with present knowledge of the evolution of Og., One can determine values compatible with present knowledge of the evolution of $\Omega_{\rm HI}$.1157 Under the assumption ofEqn., Under the assumption ofEqn.1158 18. the total neutral hydrogen density of Eqn., \ref{eq:assumption} the total neutral hydrogen density of Eqn.1159 17 is given by Using the cosmic density of neutral gas £3.(2) (measured relative to the present day critical density) one can estimate the value of 3 by assuming that the fraction of HE and He is constant between z=0 and z=1., \ref{eq:HIdensity} is given by Using the cosmic density of neutral gas $\Omega_{\rm g}(z)$ (measured relative to the present day critical density) one can estimate the value of $\beta$ by assuming that the fraction of HI and He is constant between $z=0$ and $z=1$.1160" By taking the ratio of measured neutral gas cosmic densities in the local Universe. O.(z20)23.5.10. and from camped Lyman-o systenis at higher-redshift (Prochaskactal.2005:RaoandTurnshek 2000).. OQ,(z21)=1.0.10 * ve can estimate J—2, which is also compatible with other works (Lahetal.2007).."," By taking the ratio of measured neutral gas cosmic densities in the local Universe, $\Omega_{\rm g}(z\approx 0 )=3.5 \times 10^{-4}$, and from damped $\alpha$ systems at higher-redshift \citep{Prochaska2005,RaoTurnshek}, $\Omega_{\rm g}(z\approx 1)=1.0\times 10^{-3}$ , we can estimate $\beta\sim 2$, which is also compatible with other works \citep{Lah:2007nk}."1161 The elleets of including this mass evolution for are shown in Fig., The effects of including this mass evolution for are shown in Fig.1162" 6 for the 100-beam FAST with £a,=6005."," \ref{fig:evonumber} for the 100-beam FAST with $t_{\rm obs}=600\,s$."1163 It is clear that evolution leads to an increase in the number of galaxies found. per day. which ave 14420. 19070. 25190. 321780 [or 3=0.1.2.3 and the median redshift increases with S.," It is clear that evolution leads to an increase in the number of galaxies found per day which are 14420, 19070, 25190, 32780 for $\beta=0,1,2,3$ and the median redshift increases with $\beta$."1164 Essentially. we have show that one can take the current prediction of ~10* galaxies found by the L00-beam FAST in 2 vears of observation is a conservative lower bound.," Essentially, we have show that one can take the current prediction of $\sim 10^{7}$ galaxies found by the 100-beam FAST in 2 years of observation is a conservative lower bound."1165 Η we take seriously the value of 3=2 then one would expec to find nearly twice as many galaxies as in the case of no evolution., If we take seriously the value of $\beta=2$ then one would expect to find nearly twice as many galaxies as in the case of no evolution.1166 Prior to embarking on the full. all-sky galaxy recishif survey. it would be sensible to perform. an exploratory. deeper survey of a much smaller area. with the objective of measuring the evolution of the LL mass function.," Prior to embarking on the full, all-sky galaxy redshift survey, it would be sensible to perform an exploratory, deeper survey of a much smaller area, with the objective of measuring the evolution of the HI mass function."1167 As wel as being of legitimate scientific interest in its own ri¢h such a survey would. allow one to investigate the optima depth of the main redshift survey ancl pin-down some of the questions raised in this section., As well as being of legitimate scientific interest in its own right such a survey would allow one to investigate the optimal depth of the main redshift survey and pin-down some of the questions raised in this section.1168 In Fig., In Fig.1169" 7 we show the redshift cüstribution of galaxies that one would expect. for an HIE survey with /,4,;—6000s lasting 30d using the Ξ0 mass function and bins As= 0.1."," \ref{fig:deepsurvey} we show the redshift distribution of galaxies that one would expect for an HI survey with $t_{\rm obs}=6000s$ lasting $30\,d$ using the $z=0$ mass function and bins $\Delta z=0.1$ ."1170 One would expected to find around 42000 galaxies with (2)m0.8 using mp=19 and 222000 for ni=100., One would expected to find around 42000 galaxies with $\langle z\rangle\approx 0.3$ using $n_{\rm B}=19$ and 222000 for $n_{\rm B}=100$.1171 In order to estimate accurately the parameters of the mass function. we estimate that one would require 1000 galaxies per bin.," In order to estimate accurately the parameters of the mass function, we estimate that one would require $\sim 1000$ galaxies per bin."1172 Pherefore. we believe that it should be possible to determine the HE mass function out to z0.6 for mi=19 and to zz1 np=100 with such a month-long survey.," Therefore, we believe that it should be possible to determine the HI mass function out to $z\approx 0.6$ for $n_{\rm B}=19$ and to $z\approx 1$ $n_{\rm B}=100$ with such a month-long survey."1173 We now use the precicted galaxy number counts. assuming no evolution of the kind. described in the last section. to estimate the errors ofthe galaxy power spectrum at z= 0.15.," We now use the predicted galaxy number counts, assuming no evolution of the kind described in the last section, to estimate the errors ofthe galaxy power spectrum at $z=\langle z\rangle\approx 0.15$ ."1174 Ph.z) is related to the power spectrum Pk.0) by," $P(k,z)$ is related to the power spectrum $P(k,0)$ by"1175stratification of white chyvarls is fixed in prior stages of their evolution. the evolutionary history of progenitor stars is of utmost importance in the context of white dwarf asteroseismology.,"stratification of white dwarfs is fixed in prior stages of their evolution, the evolutionary history of progenitor stars is of utmost importance in the context of white dwarf asteroseismology."1176 Our asteroscismological approach. while being physically sounding. is bv far much more computationally demancing than other approaches in which simplified. models are used.," Our asteroseismological approach, while being physically sounding, is by far much more computationally demanding than other approaches in which simplified models are used."1177 As a result. our. approach severely limits the exploration of the parameter space of he moclels.," As a result, our approach severely limits the exploration of the parameter space of the models."1178 Indeed. for the case of DA white cwarls. we iive only two parameters which we are able to vary in a consistent wav: the stellar mass (419) ancl the cllective emperature (Ziae).," Indeed, for the case of DA white dwarfs, we have only two parameters which we are able to vary in a consistent way: the stellar mass $M_*$ ) and the effective temperature $T_{\rm eff}$ )."1179 Instead. the thickness of the LL envelope (Mg). the content of Le (Mg). the shape of the C- chemical structure at the core (including the precise oportions of central O and ©). and the thickness of the chemical transition regions are fixed by the evolutionary uistory of progenitor stars.," Instead, the thickness of the H envelope $M_{\rm1180 H}$ ), the content of He $M_{\rm He}$ ), the shape of the C-O chemical structure at the core (including the precise proportions of central O and C), and the thickness of the chemical transition regions are fixed by the evolutionary history of progenitor stars."1181 Therefore. to push on the limits of our asteroseismological exploration. it would be desirable o change some additional parameters besides the stellar mass and elfective temperature of our. DA mocels.," Therefore, to push on the limits of our asteroseismological exploration, it would be desirable to change some additional parameters besides the stellar mass and effective temperature of our DA models."1182 In this work. we have chosen to vary the thickness of the LI envelope. because of the uncertainties in the mass loss rates.," In this work, we have chosen to vary the thickness of the H envelope, because of the uncertainties in the mass loss rates."1183 According to full evolutionary computations (Althaus et al., According to full evolutionary computations (Althaus et al.1184" 2000b). the maximum II envelope mass expected in a white dwarf depends on the stellar mass and ranges from ΑΗΛΙ~2.4.101 (for Ad,= 0.52534.) t6 8.510."" (for Al,—0STSAM \(see the first row of Table 13)."," 2010b), the maximum H envelope mass expected in a white dwarf depends on the stellar mass and ranges from $M_{\rm H}/M_* \sim 2.4 \times 10^{-4}$ (for $M_*= 0.525 M_{\odot}$ ) to $8.5 \times 10^{-6}$ (for $M_*= 0.878 M_{\odot}$ )(see the first row of Table \ref{table1}) )."1185 Our decision for changing this parameter is due to several reasons: first. there are compelling theoretical reasons to believe that the II-content of DA white cdwarfs might depend on the details of their previous evolution.," Our decision for changing this parameter is due to several reasons: first, there are compelling theoretical reasons to believe that the H-content of DA white dwarfs might depend on the details of their previous evolution."1186 On the contrary. the He content or the inner C-O chemical profiles are not expected to vary significantIy due to the details ofthe previous evolutionary uistory (with the exception of a possible merger origin for he white ενας).," On the contrary, the He content or the inner C-O chemical profiles are not expected to vary significantly due to the details of the previous evolutionary history (with the exception of a possible merger origin for the white dwarfs)."1187 Indeed. the total LE content remaining in some DA white dwarks could be several orders of magnituce ower than that. predicted. by our standard. treatment. of oogenitor evolution.," Indeed, the total H content remaining in some DA white dwarfs could be several orders of magnitude lower than that predicted by our standard treatment of progenitor evolution."1188 For instance. Althaus ct al. (," For instance, Althaus et al. ("11892005b) ave found that Ady becomes considerably reduced. if the »ogenitor experiences a late thermal pulse episode (LLL) shortly after the departure from the thermallv-pulsing AGB yhase.,2005b) have found that $M_{\rm H}$ becomes considerably reduced if the progenitor experiences a late thermal pulse episode (LTP) shortly after the departure from the thermally-pulsing AGB phase.1190 In this sense. ‘Tremblay Bergeron (2008) show hat the increase in the ratio of Le-rich to H-rich. white cdwarfs can be understood on the basis that a fraction. of DA white cwarls above Zip710000 Ix are characterize » a broad range of L-laver thickness.," In this sense, Tremblay Bergeron (2008) show that the increase in the ratio of He-rich to H-rich white dwarfs can be understood on the basis that a fraction of DA white dwarfs above $T_{\rm eff} \approx 10\,000$ K are characterized by a broad range of H-layer thickness."1191 Second. the precise ocation of the Πο transition region (and the value of Aly) strongly alfects the structure of the adiabatie perioc spectrum in a DA white dwarf (Bradley. 1996).," Second, the precise location of the He/H transition region (and the value of $M_{\rm H}$ ) strongly affects the structure of the adiabatic period spectrum in a DA white dwarf (Bradley 1996)."1192 Finally. Mg is the structural parameter that can be more easily mocifiec in our models without removing relevant. features predictec by the complete progenitor evolution.," Finally, $M_{\rm H}$ is the structural parameter that can be more easily modified in our models without removing relevant features predicted by the complete progenitor evolution."1193 In order to get dillerent. thicknesses of the LE envelope. we have followed a simple. recipe.," In order to get different thicknesses of the H envelope, we have followed a simple recipe."1194 For cach sequence characterized by a given stellar mass and a thick value of Ag. as predicted by the full computation of the pre-white dwarf evolution. (second. row of Table 13). we have simply replaced !HI by tHe at the basis of the LE envelope.," For each sequence characterized by a given stellar mass and a thick value of $M_{\rm H}$ , as predicted by the full computation of the pre-white dwarf evolution (second row of Table \ref{table1}) ), we have simply replaced $^{1}$ H by $^{4}$ He at the basis of the H envelope."1195 This is done at very high effective temperatures (>70000 Ix). in such a wav that the unphysical transitory elfectsx associated to this procedure end. much long before the models reach the stage of pulsating DA white dwarfs.," This is done at very high effective temperatures $\gtrsim 70\, 000$ K), in such a way that the unphysical transitory effects associated to this procedure end much long before the models reach the stage of pulsating DA white dwarfs."1196 After ourfioc procedure to change the thickness of the LE envelope. we allow time-dependent element. dilfusion to operate while the models cool down until they reach the effective temperatures characterizing the DAW instability strip.," After our procedure to change the thickness of the H envelope, we allow time-dependent element diffusion to operate while the models cool down until they reach the effective temperatures characterizing the DAV instability strip."1197 Dilfusion leacis to very smooth chemical. profiles at the Llesll chemical transition regions., Diffusion leads to very smooth chemical profiles at the He/H chemical transition regions.1198 The resulting values of the LE content for the different envelopes are shown in Table 1.. and a graphical representation of the basic grid of models employed in this work is displaved in Fig. 1l.," The resulting values of the H content for the different envelopes are shown in Table \ref{table1}, and a graphical representation of the basic grid of models employed in this work is displayed in Fig. \ref{grid_mh}."1199 In this figure. the canonical values of Ag predicted. by stellar evolution are connected with a thick (orange) linc.," In this figure, the canonical values of $M_{\rm H}$ predicted by stellar evolution are connected with a thick (orange) line."1200 Obviously. bevond the availability of the models of this coarse grid. we have the capability to generate additional DA white chart evolutionary sequences with arbitrary values of Mg for each stellar mass in order to refine the model grid.," Obviously, beyond the availability of the models of this coarse grid, we have the capability to generate additional DA white dwarf evolutionary sequences with arbitrary values of $M_{\rm H}$ for each stellar mass in order to refine the model grid."

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