ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2" This is unsurprising, given the obvious relationship between temperature and thermal broadening and between inclination and rotational broadening."," This is unsurprising, given the obvious relationship between temperature and thermal broadening and between inclination and rotational broadening."3" The optically thick CO(3-2) line responds only weakly to variations in density, and the outer radius and position angle of emission should intuitively be unrelated to line broadening, hence the independence of turbulent linewidth from 61, R., and PA."," The optically thick CO(3-2) line responds only weakly to variations in density, and the outer radius and position angle of emission should intuitively be unrelated to line broadening, hence the independence of turbulent linewidth from $c_1$, $R_c$, and PA."4 For the D'Alessio et al., For the D'Alessio et al.5" models, inclination and CO abundance (i and Xco) have the strongest relationships with turbulent linewidth."," models, inclination and CO abundance $i$ and $X_{CO}$ ) have the strongest relationships with turbulent linewidth."6" The contribution of the CO abundance in this case can be understood as a thermal broadening effect: because of the vertical temperature gradient (see Figure 6)), the CO abundance controls the location of the 7—1 surface and therefore the apparent temperature of the CO(3-2) line emission."," The contribution of the CO abundance in this case can be understood as a thermal broadening effect: because of the vertical temperature gradient (see Figure \ref{fig:temp_dens}) ), the CO abundance controls the location of the $\tau$ =1 surface and therefore the apparent temperature of the CO(3-2) line emission."7" To characterize the effects of these variables on the observable properties of the CO(3-2) emission, we investigate their influence on a toy model of optically thick line emission."," To characterize the effects of these variables on the observable properties of the CO(3-2) emission, we investigate their influence on a toy model of optically thick line emission."8" We assume a power-law temperature distribution for a geometrically flat, optically thick, azimuthally symmetric circumstellar disk."," We assume a power-law temperature distribution for a geometrically flat, optically thick, azimuthally symmetric circumstellar disk."9" In the Rayleigh-Jeans approximation, the brightness of the line at a given frequency will be directly proportional to the temperature."," In the Rayleigh-Jeans approximation, the brightness of the line at a given frequency will be directly proportional to the temperature."10" We include two sources of line broadening, thermal and turbulent, implemented by the relationship Av(r)= where Av is the total linewidth, € is the turbulent ./2keT(r)/m+linewidth,£2, and the thermal linewidth is where kg is Boltzmann’s constant, T is the local \/2kgT/mtemperature in the disk, and m is the average mass per particle."," We include two sources of line broadening, thermal and turbulent, implemented by the relationship $\Delta v(r) = \sqrt{2 k_B T(r)/m + \xi^2}$ , where $\Delta v$ is the total linewidth, $\xi$ is the turbulent linewidth, and the thermal linewidth is $\sqrt{2 k_B T/m}$ where $k_B$ is Boltzmann's constant, $T$ is the local temperature in the disk, and $m$ is the average mass per particle."11 Rotational broadening isimplicitly included in the assumed Keplerian rotation, Rotational broadening isimplicitly included in the assumed Keplerian rotation12is to say. the first) ane second. bodies do not collide with cach other or any other bodies than the zeroth.,"is to say, the first and second bodies do not collide with each other or any other bodies than the zeroth."13 In practice. simulations redefine the relative coordinates whenever necessary. according to who is close to whom.)," In practice, simulations redefine the relative coordinates whenever necessary, according to who is close to whom.)"14 For the equations of motion we need. derivatives with respect to quaternion components., For the equations of motion we need derivatives with respect to quaternion components.15" First we have Vo,Q7= 391.", First we have $\nablaQ Q_1^2 = 2\Q_1$ .16" Slightly more subtle is Vg,reQjA]=A il A is independent of Qj.", Slightly more subtle is $\nablaQ \tr{\Q_1^*\A}=\A$ if $\A$ is independent of $\Q_1$.17 Using this last identity. together with the definition (31)) of LL. we derive Hamilton's equations are then and similarly for P».Q..," Using this last identity, together with the definition \ref{dotasPi}) ) of $\Pi$, we derive Hamilton's equations are then and similarly for $\P_2,\Q_2$."18 In dvnamical astronomy the IWS transformation is profound. but may appear mysterious.," In dynamical astronomy the KS transformation is profound, but may appear mysterious."19 This paper attempts to make ib less mysterious. ancl hopefully. therefore more. useful. bv explaining it in three-dimensional ecometric terms.," This paper attempts to make it less mysterious, and hopefully therefore more useful, by explaining it in three-dimensional geometric terms."20 There are several possible directions in which the Ks transformation may turn out to be useful., There are several possible directions in which the KS transformation may turn out to be useful.21 First. one can imagine new orbit integrators specialized to nearlv-Ixeplerian problems.," First, one can imagine new orbit integrators specialized to nearly-Keplerian problems."22 Work on dense stellar systems with near collisions has. alreacky been mentioned. (for.re-viewsseethebooks ??)..," Work on dense stellar systems with near collisions has already been mentioned \citep[for reviews see the23books][]{2003gnbs.book.....A,2003gmbp.book.....H}."24 In the planetary regime. which cilfers from the dense-stellar case in having few bodies bu many more orbital times. time transformations reminiscen of (28)) used. for WS regularization have proved: usefu for highly eccentric orbits (77).. while some integration algorithms (77) apply the time transformation (28)) implicitly.," In the planetary regime, which differs from the dense-stellar case in having few bodies but many more orbital times, time transformations reminiscent of \ref{kepeq}) ) used for KS regularization have proved useful for highly eccentric orbits \citep{1997CeMDA..67..145M,2002CeMDA..84..331E}, while some integration algorithms \citep{1999CeMDA..74..287M,1999AJ....118.2532P} apply the time transformation \ref{kepeq}) ) implicitly."25 C'ould the KS transformation itself be exploitec here?, Could the KS transformation itself be exploited here?26 ? has some further ideas., \cite{2005AJ....129.2496F} has some further ideas.27 Second. it is conceivable that KS variables: coul simply perturbation theory.," Second, it is conceivable that KS variables could simplify perturbation theory."28 Perturbation theory in Classical celestial mechanics (seeforexample?) is algebraically frighteningly complicated. basically because the natural variables for the unperturbecl ancl perturbed parts (being the Weplerian action-angles and real-space coordinate) are related. through an implicit equation.," Perturbation theory in classical celestial mechanics \citep[see for example][]{2000ssd..book.....M} is algebraically frighteningly complicated, basically because the natural variables for the unperturbed and perturbed parts (being the Keplerian action-angles and real-space coordinate) are related through an implicit equation."29= On the other hand. the action-angles of the IxXS-transformed Ixepler problem are explicitly related to space coordinatesthe implicit equation is transferred. to the time variable.," On the other hand, the action-angles of the KS-transformed Kepler problem are explicitly related to space coordinates---the implicit equation is transferred to the time variable."30 Could some major simplication be achieved. through Ks variables?, Could some major simplication be achieved through KS variables?31 Some progress has been mace by ?.., Some progress has been made by \cite{2006NewA...11..366V}.32 Third. the INS transformation might provide new insight into analogous quantum problem.," Third, the KS transformation might provide new insight into analogous quantum problem."33 ?? derive the symmetry groups of the bound and unbound Coulomb problem.," \cite{RevModPhys.38.330,RevModPhys.38.346} derive the symmetry groups of the bound and unbound Coulomb problem."34 These turn out to be the same four-dimensional svmimetries as in IxS theory., These turn out to be the same four-dimensional symmetries as in KS theory.35 Is the KS transformation implicit in that work?, Is the KS transformation implicit in that work?36 Lam grateful to thank Seppo Alikkola for introducing me to WS theory. and to Marcel Zemp. Scott Tremaine. and the referee. JOrre Waldvogel. for suggesting improvements in the nianuscript.," I am grateful to thank Seppo Mikkola for introducing me to KS theory, and to Marcel Zemp, Scott Tremaine, and the referee, Jörrg Waldvogel, for suggesting improvements in the manuscript."37"where e is defined in eq.(1). and (he superscripts Ο ancl S. denote observed image ancl source. respectively,","where $\mathbf{e}$ is defined in eq.(1), and the superscripts 'O' and 'S' denote observed image and source, respectively."38" Then (he noisy convergence &,, including the contamination Iron source elliplicilies follows Considering smoothed quantities. we have (e.g.. van Waerbeke 2000) ancl where X. P. and Avy(0) are the smoothed οἱο). 5 and &,. respectively. WW(4) is the smoothing function. and n», and .N, are. respectively. the surface number density and the nunmber of source galaxies in (he field."," Then the noisy convergence $\kappa_n$ including the contamination from source ellipticities follows Considering smoothed quantities, we have (e.g., van Waerbeke 2000) and where $\mathbf{\Sigma}^{(O)}$ , $\mathbf{\Gamma}$, and $K_N(\vec \theta)$ are the smoothed $\mathbf{e^{(O)}}$, $\gamma$ and $\kappa_n$ , respectively, $W(\vec \theta)$ is the smoothing function, and $n_g$ and $N_g$ are, respectively, the surface number density and the number of source galaxies in the field."39 The noise part of A due to the intrinsic ellipticities is then where HV(&) is the Fourier transformation of the smoothing function with the form Following van Waerbeke (2000). the correlation of NV(4) is caleulated by averaging over both the ellipticities and the positions of source galaxies.," The noise part of $K_N$ due to the intrinsic ellipticities is then where $\tilde W(\vec k)$ is the Fourier transformation of the smoothing function with the form Following van Waerbeke (2000), the correlation of $N(\vec \theta)$ is calculated by averaging over both the ellipticities and the positions of source galaxies."40" Without intrinsic alignments. the correlation of N(@) arises only [rom the smoothing operations. and by ignoring the sampling of source galaxies. we have (van Waerbeke 2000) where o, is the intrinsic dispersion of e). and the factor (27)? comes in to be in accord wilh the definition of W(&) in eq. ("," Without intrinsic alignments, the correlation of $N(\vec \theta)$ arises only from the smoothing operations, and by ignoring the non-uniform sampling of source galaxies, we have (van Waerbeke 2000) where $\sigma_{\epsilon}$ is the intrinsic dispersion of $\mathbf{e}^{(S)}$, and the factor $(2\pi)^2$ comes in to be in accord with the definition of $\tilde W(\vec k)$ in eq. ("4115).,15).42 Including the alignment. theoperation by averaging over (he ellipticiüesof source ealaxies. denoted by ;1 following van Waerbeke (2000). is," Including the alignment, theoperation by averaging over the ellipticitiesof source galaxies, denoted by ${\it {A}}$ following van Waerbeke (2000), is"43 , 44(he solar convection zone than it does al the Earth's tropopause.,the solar convection zone than it does at the Earth's tropopause.45 These differences can have prolound effects on the role of overshoot and. hence. on the scale. frequency. and amplitude of (he waves generated and. so. on the angular momentum transport bv these waves.," These differences can have profound effects on the role of overshoot and, hence, on the scale, frequency and amplitude of the waves generated and, so, on the angular momentum transport by these waves."46 Numerical simulations (Wedi Smolarkiewicz 2005) of the Phumb-MeEswan laboratory experiment attempting to reproduce (he QDO. have shown (hat the (wpe and period of an oscillation in the differential rotation profile depend sensitively on the lorcing.," Numerical simulations (Wedi Smolarkiewicz 2005) of the Plumb-McEwan laboratory experiment attempting to reproduce the QBO, have shown that the type and period of an oscillation in the differential rotation profile depend sensitively on the forcing."47 In particular. ib is found that random forcing rarely produces a periodic oscillation.," In particular, it is found that random forcing rarely produces a periodic oscillation."48 Given the turbulent nature of the sun. it is likely that the forcing is fairly random.," Given the turbulent nature of the sun, it is likely that the forcing is fairly random."49 For reasons stated above. it is unlikely (hat there is à QBO-like oscillation associated with the solar tachcocline.," For reasons stated above, it is unlikely that there is a QBO-like oscillation associated with the solar tachcocline."50 Lowever. QDO-like oscillations may occur in stars will radiative envelopes because of the inelliciency of overshoot into an overlving stable region. and because ol more similar geometry.," However, QBO-like oscillations may occur in stars with radiative envelopes because of the inefficiency of overshoot into an overlying stable region, and because of more similar geometry."51 We have presented self-consistent numerical simulations of convective overshoot ancl eravily wave generation and (he angular momentum (ransport by these processes in a 2D model of the dynamics in the solar equatorial plane., We have presented self-consistent numerical simulations of convective overshoot and gravity wave generation and the angular momentum transport by these processes in a 2D model of the dynamics in the solar equatorial plane.52 We find (hat angular velocity variations in (he tachocline are driven by angular momentum t(rausported by overshooting plumes rather (han nonlinear interaction of low amplitude waves., We find that angular velocity variations in the tachocline are driven by angular momentum transported by overshooting plumes rather than nonlinear interaction of low amplitude waves.53 These overshooting plumes are strongly nonlinear disturbances. which can not be accurately represented as an increased fIux of linear waves.," These overshooting plumes are strongly nonlinear disturbances, which can not be accurately represented as an increased flux of linear waves."54 We observe a semi-periodic oscillation in amplitude. but not in direction. of (he mean flow in the tachocline because of an asvimnietry in the driving of prograde and retrograde motions.," We observe a semi-periodic oscillation in amplitude, but not in direction, of the mean flow in the tachocline because of an asymmetry in the driving of prograde and retrograde motions."55 Since we find that linear gravity waves are not dominant in the tachocline it is unlikely Chat thev are responsible for the 1.3 vear oscillation or the 11 vear solar cvele., Since we find that linear gravity waves are not dominant in the tachocline it is unlikely that they are responsible for the 1.3 year oscillation or the 11 year solar cycle.56 It is no surprise (hat overshooting motions plav a dominant role in the tachocline aud we expect these results will persist in three-dimensions., It is no surprise that overshooting motions play a dominant role in the tachocline and we expect these results will persist in three-dimensions.57 In the deep radiative interior the continual deposition of angular momentum by the Yonlinear interaction of gravity waves produces a radiallv banded differential rotation., In the deep radiative interior the continual deposition of angular momentum by the nonlinear interaction of gravity waves produces a radially banded differential rotation.58 However. il remains {ο be seen whether this pattern persists in 3D. considering its verv low amplitude.," However, it remains to be seen whether this pattern persists in 3D, considering it's very low amplitude."59 In the models core. the amplitude of the differential rotation (ie.. angular velocitv) is arger. about (vo orders of magnitude larger than that in the bulk of the radiative region and similar to the magnitude within the convection zone.," In the model's core, the amplitude of the differential rotation (i.e., angular velocity) is larger, about two orders of magnitude larger than that in the bulk of the radiative region and similar to the magnitude within the convection zone."60 We observe retrograde motion in (he core reversing (o prograde motion in step with the counter-reversal (prograde (o retrograde) at the tachocline., We observe retrograde motion in the core reversing to prograde motion in step with the counter-reversal (prograde to retrograde) at the tachocline.61 When (here is predominantly prograde flow at the base of, When there is predominantly prograde flow at the base of62 1 (IXomossactal.2003.seeColpi& ~ (ATaness," $\lsim 1$ \citep[][see Colpi \& Dotti 2009 for63 a recent review]{komossa2003}."64etal.2001:Rodriguez2006).," $\sim$ \citep{maness2004,rodriguez2006}."65 z7 (2=0.055) (seeValtonenetal.2008.andreferencestherein).. zz12 (Bee," $\approx 7$ $z = 0.055$ \citep[see][and references66 therein]{valtonen2008}."67chnan.Dlaudford&Rees1980)., $\approx 12$ \citep{begelman1980}.68 stripped by the eravitational poteutial of the companion. resulting iu peculiar flux ratios between DLs with different ionization potential (Montuorietal.2010).," stripped by the gravitational potential of the companion, resulting in peculiar flux ratios between BLs with different ionization potential \citep{montuori2010}."69. This spectroscopic approach does not suffer any aueular resolution limitations: Actually. the closer (aud more massive} the binary is. the more shifted/deforined the DLs are.," This spectroscopic approach does not suffer any angular resolution limitations: Actually, the closer (and more massive) the binary is, the more shifted/deformed the BLs are."70" Thaulss to the existence of large spectroscopic survevs, such as the Sloan Digital Skv Survey (SDSS). a laree region of the skv can be probed."," Thanks to the existence of large spectroscopic surveys, such as the Sloan Digital Sky Survey (SDSS), a large region of the sky can be probed."71 Up to date five spectroscopically identified candidates have Όσοι preseuted: J0927|2913 (Ixoiossa.Zhou&Lu2008:Dos-dauovic.Eracleous&Sigurdsson2009:Dottietal. 2009). J1536|0111. CBorosoun&Lauer2009).. J1050|3156 (Shieldsetal.2009)... 1€|22.25 (J1000|2233per.Decarlietal. 201053... and JO9320318 (Barrowsctal. 2011)..," Up to date five spectroscopically identified candidates have been presented: J0927+2943 \citep{komossa08,bogdanovic09,dotti09}, J1536+0441 \citep{boroson09}, J1050+3456 \citep{shields09}, 4C+22.25 \citep[J1000+2233 in this paper,][]{decarli_4c2225}, and J0932+0318 \citep{barrows11}. ."72 Such à simall umber of objects| is mareiually compatible with the theoretically predicted. umber of sub-parsec BUBs at τς0.7 (510.eiventhemergerrateandobservabilitv:seeVolouterietal. 2009).," Such a small number of objects is marginally compatible with the theoretically predicted number of sub-parsec BHBs at $z \lsim 0.7$ \citep[5--10, given the merger rate and73under reasonable assumptions on the binary lifetime and observability; 74see][]{volonteri09}."75. The spectroscopic approach las an obvious drawback: a peculiar spectrum does not euarantee the presence of a DIID iu the uucleus of the host., The spectroscopic approach has an obvious drawback: a peculiar spectrum does not guarantee the presence of a BHB in the nucleus of the host.76 As an example. an nnobscirecdk BIB with both DIIs active could resaiid5 the spectrum of a double peaked. emitter (see.e.g..Er-acleous&Παρα 1991).. where broad doublepeaked lines are cluitted because of the almost οσοon. disklike structure of the DL region ofa single BIT.," As an example, an unobscured BHB with both BHs active could resamble the spectrum of a double peaked emitter \citep[see,77 e.g.,][]{eracleous1994}, where broad double–peaked lines are emitted because of the almost edge–on, disk–like structure of the BL region of a single BH."78 A binary with a single accreting DIT would show a single shifted BL., A binary with a single accreting BH would show a single shifted BL.79 It the shift corresponds to a relatively small velocity along the line of sightCZ1000 kun s +). the same signature could be emitted by a remnant ofa binary coalescence. recoiling because of anisotropicgravitational wave cussion (e.g.," If the shift corresponds to a relatively small velocity along the line of sight$\lsim 4\,000$ km $^{-1}$ ), the same signature could be emitted by a remnant of a binary coalescence, recoiling because of anisotropicgravitational wave emission \citep[e.g.][]{komossa08}"80QPOs possidv. observed from £U 155030 (Zhaue 11998) a signature of the preseuce of the innermost stable circular orbit 15600). a crucial prediction of stroug-eravitv general relativity.,"QPOs possibly observed from 4U 1820–30 (Zhang 1998) a signature of the presence of the innermost stable circular orbit (ISCO), a crucial prediction of strong-gravity general relativity."81 σος only iu BFAIs cau one infer a Hel gravitational mass M>ολ. for this source. which constraius strouglv the equation of state of the high-density matter in fre core of neutron stars.," Hence only in BFMs can one infer a high gravitational mass $M>2.1M_\odot$ for this source, which constrains strongly the equation of state of the high-density matter in the core of neutron stars."82 The chawine separation frequency observed iu several sources provided part of the motivation for t16 development of other models of the silohertz QPOs. iu utieular the relativisic precession models (e.g.. Stella Vietzii 1998).," The changing separation frequency observed in several sources provided part of the motivation for the development of other models of the kilohertz QPOs, in particular the relativistic precession models (e.g., Stella Vietri 1998)."83 In these nodels the close match between t16 separation frequency aud the spin frequency interred from burst brielitucss oscilations is a coincideice. but they do precict the qualitative effect ofa searation frequency hat drops with increasing kilohertz QPO YOequencev.," In these models the close match between the separation frequency and the spin frequency inferred from burst brightness oscillations is a coincidence, but they do predict the qualitative effect of a separation frequency that drops with increasing kilohertz QPO frequency."84 Tere we discuss the beat-frequeACY nocclin helt of hese new developeits ancl contrast it with alternate pictures., Here we discuss the beat-frequency model in light of these new developments and contrast it with alternate pictures.85 Ins 2 we describe tιο observaional trends that notivated the development of bea-frequeicv models., In 2 we describe the observational trends that motivated the development of beat-frequency models.86 We then οaborate oi these uodels; iu particular the souic-poiut beat-frequency model.," We then elaborate on these models, in particular the sonic-point beat-frequency model."87 In 3Pa we ¢iscuss the evicence for a changine difference yequenc wodmn the four sources Sco X-1. IU IT 172831. aud. IU 1735|ο," In 3 we discuss the evidence for a changing difference frequency in the four sources Sco X-1, 4U 1608--52, 4U 1728–34, and 4U 1735–44."88" Wes row that an aspect of the sonic-poiut veat-frequency iiocdel. iucluded iu the ανiuues but originally onuittec from the yequency estimates, naturally accomunodates the changing difference frecποιον axd can quautitative vt he data."," We show that an aspect of the sonic-point beat-frequency model, included in the dynamics but originally omitted from the frequency estimates, naturally accommodates the changing difference frequency and can quantitatively fit the data."89 Finally. in Lie contrast some of he predictions of he beat-frequency uodoel with the predictions of he relativistic precession model. aud cliscuss analysis that ασ! be doue with ctirent data to help discriminate vetween the two iteroxetatious.," Finally, in 4 we contrast some of the predictions of the beat-frequency model with the predictions of the relativistic precession model, and discuss analysis that might be done with current data to help discriminate between the two interpretations."90 We also explore fje inurpact of future observations. oth with the u»voniues eeneration of Ligh spectral resolution satellites (such as Chandra and NMBM) aid with louger-terii projecs such as Constellation-N. aud a wpothetical ligarea follow-on toXTE.," We also explore the impact of future observations, both with the upcoming generation of high spectral resolution satellites (such as Chandra and XMM) and with longer-term projects such as Constellation-X and a hypothetical high-area follow-on to."91 As discussed iu the iitroductiou. soon after the discovery of kiloverte QPOs it was established tha these oscilations have (1) higoh frequency. (2) high amplitide. aud (3) high colereicc. àid that there are always two or fewer kiohertz QPOs iu a eiveu oer density spectrmu.," As discussed in the introduction, soon after the discovery of kilohertz QPOs it was established that these oscillations have (1) high frequency, (2) high amplitude, and (3) high coherence, and that there are always two or fewer kilohertz QPOs in a given power density spectrum."92 T1 addition. the separation frequency a»peared Consistct with constant mi many solrees and close to the spin frequeicy interred frou burst xiehtuess oscillatious iu he four sources where this coud be tested.," In addition, the separation frequency appeared consistent with constant in many sources and close to the spin frequency inferred from burst brightness oscillations in the four sources where this could be tested."93 Tt is LOW shown (see 23) tha in several. aud perhaps all. sources. tιο separation frequency is coustant (although 1 is still close to the interred spin frequency). ai in fact decreases systenati‘ally with increasing lower peak frequency.," It is now known (see 3) that in several, and perhaps all, sources, the separation frequency is constant (although it is still close to the inferred spin frequency), and in fact decreases systematically with increasing lower peak frequency."94 In 3 owe diseuss row this new result nay be interpreted within the beat-frequeney model., In 3 we discuss how this new result may be interpreted within the beat-frequency model.95 The Hel freqcacy indicates that the source of the brightucss oscillations is close to the neutron star., The high frequency indicates that the source of the brightness oscillations is close to the neutron star.96 A natural candidate for these oscillations is the orbital yequency at sole special radius., A natural candidate for these oscillations is the orbital frequency at some special radius.97 Caven that the burst oscillation frequency is most convincingly iutermeted as the stellar spin frequeney or its first overtone (sec. e.g.. Strolunaver AMarkwardt 1999). the close match of the separation frequency," Given that the burst oscillation frequency is most convincingly interpreted as the stellar spin frequency or its first overtone (see, e.g., Strohmayer Markwardt 1999), the close match of the separation frequency"98Radio galaxies. with linear sizes reaching up to several neeaparsecs. are possibly the largest individual objects in the Universe.,"Radio galaxies, with linear sizes reaching up to several megaparsecs, are possibly the largest individual objects in the Universe."99 It is widely accepted that thev originate roni highly energetic unou-theriial processes occurring iun he nucleus of the so-called active galaxies (Blandford Rees 197E: Rees 1978))., It is widely accepted that they originate from highly energetic non-thermal processes occurring in the nucleus of the so-called active galaxies (Blandford Rees \cite{blandford}; Rees \cite{rees1}) ).100 According to the standard model of active galactic nuclei. (ACNs). a super-massive black role with a nass between 109 and ΤοAL. resides iu he ceuter of the active galaxy. powered by an accretion disk STLPOUMLICed by a torus formed by eas aud dust.," According to the standard model of active galactic nuclei (AGNs), a super-massive black hole with a mass between $10^6$ and $10^9 M\odot$, resides in the center of the active galaxy, powered by an accretion disk surrounded by a torus formed by gas and dust."101" Iu about of these Ανα, there is intense svuchrotrou radio enissiou produced iu a bipolar ottfow of relativistic particles expelled perpendicularly to the plane of the disk aud extending to distances reaching the meeaparsec scales."," In about of these AGNs, there is intense synchrotron radio emission produced in a bipolar outflow of relativistic particles expelled perpendicularly to the plane of the disk and extending to distances reaching the megaparsec scales."102 The reason why an ACN prescuts or not powerful radio enüssion is matter of strong debate., The reason why an AGN presents or not powerful radio emission is matter of strong debate.103 While there is Increasing evidence about the existence of superanassive black holes in the center of AGNs. iux even at the uuclei of nou active galaxies (Macchetto. 1999:: I&onueudy. Gebhardt 2001)). if has been argued tiat the preseuce or uot of iuteuse radio emission might be due to the rotatiou velocity of the black hole (e.g. Wilson Colbert 1995: Cavaliere D'Elia 20023). or to its total mass aud the efficiency. of accretion (àcLure Duulop 2001:: Dunlop et al. 2003).," While there is increasing evidence about the existence of super-massive black holes in the center of AGNs, and even at the nuclei of non active galaxies (Macchetto \cite{macchetto}; Kormendy Gebhardt \cite{kormendy}) ), it has been argued that the presence or not of intense radio emission might be due to the rotation velocity of the black hole (e.g. Wilson Colbert \cite{wilson}; Cavaliere D'Elia \cite{cavaliere}) ), or to its total mass and the efficiency of accretion (McLure Dunlop \cite{mclure}; Dunlop et al. \cite{dunlop2}) )."104 Considering a natural evolutionary sequence of radio ealaxies. jets clmanating from the couter of activity start boring their way fwoueh the iuterstellar mediuu first. reaching the ealactic halo aud in some large cases the intergalactic uediunu.," Considering a natural evolutionary sequence of radio galaxies, jets emanating from the center of activity start boring their way through the interstellar medium first, reaching the galactic halo and in some large cases the intergalactic medium."105 Finally. the lobes of radio galaxies which have ceased them central activity expand aud disappear in the external medium.," Finally, the lobes of radio galaxies which have ceased their central activity expand and disappear in the external medium."106 Badio sources represcuting these differeut pliases of evolution are currently known adding support to this scenario: compact svuunetrie objects (CSOs: Wilkinson ct al. 1991)), Radio sources representing these different phases of evolution are currently known adding support to this scenario: compact symmetric objects (CSOs; Wilkinson et al. \cite{wilkinson}))107" are thought to be voung radio eaaxies (ος, Owsianilk Comvay 1998)). while the eianut radio galaxies (GRCGs: defined as those with a projected linear > 1 Mpc) are probably old objects at the latter stages of evolution (slovara-Chancdra Saikia 1999))."," are thought to be young radio galaxies (e.g. Owsianik Conway \cite{owsianik}) ), while the giant radio galaxies (GRGs; defined as those with a projected linear $\ge$ 1 Mpc) are probably old objects at the latter stages of evolution (Ishwara-Chandra Saikia \cite{ishwara}) )."108 Relic radio sources found in clusters of galaxies ποτ correspond to the last detectable cussion from “dead” radio galaxies (e.g. TNomissarov Cubanov 1991: Slee et al. 2001))., Relic radio sources found in clusters of galaxies might correspond to the last detectable emission from “dead” radio galaxies (e.g. Komissarov Gubanov \cite{komissarov}; Slee et al. \cite{slee}) ).109 However. the degree of influence of parameters other than the age (e.g. source power. conditions of the external mediun)iu the evolution of radio galaxies is not clear.," However, the degree of influence of parameters other than the age (e.g. source power, conditions of the external medium) in the evolution of radio galaxies is not clear."110 For example. although there is observational evidence supporting the voune source scenario for CSOs. it has also been argued that CSOs are short lived objects which never reach the size of their bie relatives (Reacdhead et al. 1991)).," For example, although there is observational evidence supporting the young source scenario for CSOs, it has also been argued that CSOs are short lived objects which never reach the size of their big relatives (Readhead et al. \cite{readhead}) )."111 At he other extreme. GRGs could be the result of normal Yaio galaxies expanding in very low density environments permitting them to reach their overwhehuing sizes.," At the other extreme, GRGs could be the result of normal radio galaxies expanding in very low density environments permitting them to reach their overwhelming sizes."112 But they could also result from very powerful core activity. or both concditious must apply for a radio galaxy to become a eiut.," But they could also result from very powerful core activity, or both conditions must apply for a radio galaxy to become a giant."113 Complicating the previously outlined. evolutionary sequence. sole radio galaxies secu. to wake up after a dormant phase of absence or much lower activity (c.g. Lara et al. 1999a)).," Complicating the previously outlined evolutionary sequence, some radio galaxies seem to wake up after a dormant phase of absence or much lower activity (e.g. Lara et al. \cite{lara}) )."114 Moreover. the presence of super-auassive objects in inanv nonactive galaxies argue im favor of activity as a short trausition period in most. if not all. (elliptical) galaxies. aud that the “menace” for future activity is present at the ceuter of every. galaxy.," Moreover, the presence of super-massive objects in many non–active galaxies argue in favor of activity as a short transition period in most, if not all, (elliptical) galaxies, and that the “menace"" for future activity is present at the center of every galaxy."115 This paper is the last of a series of three devoted to the study of a sample of large augular size radio galaxies which try to address some of these open questions., This paper is the last of a series of three devoted to the study of a sample of large angular size radio galaxies which try to address some of these open questions.116 Definition of the sample and radiomaps in one side. and tages and spectroscopic data on the other. were preseuted by Lara ct al. (2001a.. ," Definition of the sample and radiomaps in one side, and images and spectroscopic data on the other, were presented by Lara et al. \cite{paperI}, ,"117hereafter Paper I) auc (2001b.. Paper," hereafter Paper I) and \cite{paperII}, , Paper"118"boundary conditions (e.g. existence of a planetary surface), chemistry, and the properties of the condensating species, such that a direct comparison with these atmospheres is difficult.","boundary conditions (e.g. existence of a planetary surface), chemistry, and the properties of the condensating species, such that a direct comparison with these atmospheres is difficult."119" Nonetheless, effective cloud formation in these atmospheres can, for example, also significantly alter the thermal emission and reflection spectra of these objects by means of both the scattering and absorption of radiation caused by cloud particles."," Nonetheless, effective cloud formation in these atmospheres can, for example, also significantly alter the thermal emission and reflection spectra of these objects by means of both the scattering and absorption of radiation caused by cloud particles."120" In this paper, we study the influence of low- and high-level clouds on the reflectance spectra and spectral albedos of Earth-like planets in the visible and NIR wavelength range at low resolution."," In this paper, we study the influence of low- and high-level clouds on the reflectance spectra and spectral albedos of Earth-like planets in the visible and NIR wavelength range at low resolution."121 A one-dimensional (1D) steady-state radiative-convective atmospheric model is used in the model calculations., A one-dimensional (1D) steady-state radiative-convective atmospheric model is used in the model calculations.122 The model accounts for two different cloud layers (low-level water droplet and high-level ice particle clouds) and for the partial overlap of these two layers., The model accounts for two different cloud layers (low-level water droplet and high-level ice particle clouds) and for the partial overlap of these two layers.123 A more detailed model description is given in Sect. 2.., A more detailed model description is given in Sect. \ref{sec_model}.124" To verify the applicability of our modelling approach, this coupled cloud-atmosphere model is applied to the modern Earth atmosphere and its spectral appearance, as described in Sect. 3.."," To verify the applicability of our modelling approach, this coupled cloud-atmosphere model is applied to the modern Earth atmosphere and its spectral appearance, as described in Sect. \ref{sec_earth_ref}."125 The resulting reflection spectra and spectral albedos of Earth-like planets orbiting different types of central stars and their implications for the detectability of characteristic molecular signatures are presented in Sect. 4.., The resulting reflection spectra and spectral albedos of Earth-like planets orbiting different types of central stars and their implications for the detectability of characteristic molecular signatures are presented in Sect. \ref{sec_spectra}.126 The impact of clouds on the thermal IR emission spectra was investigated in detail by ? (henceforth called Paper II) for Earth-like planets around different types of stars., The impact of clouds on the thermal IR emission spectra was investigated in detail by \citet{Kitzmann2011A&A531} (henceforth called Paper II) for Earth–like planets around different types of stars.127" However, clouds also affect the reflected incident stellar radiation in the short wavelength range, from the NUV to NIR, of the planetary spectrum."," However, clouds also affect the reflected incident stellar radiation in the short wavelength range, from the NUV to NIR, of the planetary spectrum."128" We study here the influence of clouds on planetary reflection spectra at low resolution using a 1D steady-state radiative-convective atmospheric model developed to account for the radiative effects of multi-layered clouds and their impact on the surface temperature in atmospheres of Earth-like planets (see?,PaperI,fordetails).."," We study here the influence of clouds on planetary reflection spectra at low resolution using a 1D steady-state radiative-convective atmospheric model developed to account for the radiative effects of multi-layered clouds and their impact on the surface temperature in atmospheres of Earth-like planets \citep[see][Paper I, for details]{Kitzmann2010}."129" We adopt a parameterised description of two different cloud layers, low-level droplet and high-level ice particle clouds, that is derived from in-situ measurements of the respective cloud type in the atmosphere of Earth."," We adopt a parameterised description of two different cloud layers, low-level droplet and high-level ice particle clouds, that is derived from in-situ measurements of the respective cloud type in the atmosphere of Earth."130" This cloud model is included into the originally cloud-free atmospheric model of ?,, in its form developed by ? and ?.."," This cloud model is included into the originally cloud-free atmospheric model of \citet{Kasting1984}, in its form developed by \citet{Pavlov00} and \citet{Segura03}."131" To limit the number of cloud parameters, the minimal possible partial overlap of both cloud layers is assumed in the calculations unless otherwise stated."," To limit the number of cloud parameters, the minimal possible partial overlap of both cloud layers is assumed in the calculations unless otherwise stated."132" The altitude of each cloud layer is not simply fixed in height, but iteratively adjusted to match the corresponding measured Earth pressure values (low-level water cloud: 0.83bar, high-level ice cloud: 0.27 bar)."," The altitude of each cloud layer is not simply fixed in height, but iteratively adjusted to match the corresponding measured Earth pressure values (low-level water cloud: $0.83~\mathrm{bar}$, high-level ice cloud: $0.27~\mathrm{bar}$ )."133 The freezing limit of the water droplets and the limiting temperature for liquefying the ice particles determine the range of temperatures for which our method is valid (cf., The freezing limit of the water droplets and the limiting temperature for liquefying the ice particles determine the range of temperatures for which our method is valid (cf.134 PaperD)., Paper.135".. For all calculations, the same chemical composition of the atmosphere is assumed, which is chosen to represent modern Earth conditions (see?).."," For all calculations, the same chemical composition of the atmosphere is assumed, which is chosen to represent modern Earth conditions \citep[see][]{Grenfell07}."136 The treatment of the radiation transfer problem in the atmospheric model is optimised for the energy transport in Earth-like atmospheres (e.g.?).., The treatment of the radiation transfer problem in the atmospheric model is optimised for the energy transport in Earth-like atmospheres \citep[e.g.][]{Mlawer97}.137" As usual, the calculation of the radiation transport is divided into two spectral parts."," As usual, the calculation of the radiation transport is divided into two spectral parts."138" The first part treats the absorption and scattering of incident solar radiation in the short wavelength regime, the second handles the absorption and emission of thermal radiation from the planetary surface and atmosphere, including the multiple scattering caused by cloud particles."," The first part treats the absorption and scattering of incident solar radiation in the short wavelength regime, the second handles the absorption and emission of thermal radiation from the planetary surface and atmosphere, including the multiple scattering caused by cloud particles."139" The plane-parallel radiative transfer equation in the short wavelength regime is solved by applying a 6-two-stream quadrature approach (?)., which uses 38 broad spectral intervals between 0.238 um and 4.55 um with variable spectral resolution (10 <R<100)."," The plane-parallel radiative transfer equation in the short wavelength regime is solved by applying a $\delta$ -two-stream quadrature approach \citep{Toon89}, which uses 38 broad spectral intervals between 0.238 $\mu$ m and 4.55 $\mu$ m with variable spectral resolution $10 \leq$ $\leq 100$ )."140" These intervals define the spectral resolution of the planetary spectra presented here, which are obtained directly from the radiative transfer calculations of the atmospheric model."," These intervals define the spectral resolution of the planetary spectra presented here, which are obtained directly from the radiative transfer calculations of the atmospheric model."141" The molecular absorption of HzO, CO», O2, O3, and CH, is treated with four- correlated-k coefficients (cf.?).."," The molecular absorption of $_2$ O, $_2$, $_2$, $_3$, and $_4$ is treated with four-term $k$ coefficients \citep[cf.][]{Segura03}."142" Rayleigh scattering is considered for O5», No, and CO; in this wavelength region."," Rayleigh scattering is considered for $_2$, $_2$, and $_2$ in this wavelength region."143 A description of the IR radiation transport treatment and the resulting thermal emission spectra is given in Paper II., A description of the IR radiation transport treatment and the resulting thermal emission spectra is given in Paper II.144" In addition to the gas opacities, the frequency-dependent optical properties of the cloud particles were previously introduced into the radiative transfer schemes including multiple scattering (Paper I)."," In addition to the gas opacities, the frequency-dependent optical properties of the cloud particles were previously introduced into the radiative transfer schemes including multiple scattering (Paper I)."145" To account for different coverages of multi-layered clouds and their partial overlap (being non uni-dimensional quantities) in a 1D atmospheric model, we use a flux-averaging procedure, where the radiative transfer problem is solved for every distinct cloud configuration separately."," To account for different coverages of multi-layered clouds and their partial overlap (being non uni-dimensional quantities) in a 1D atmospheric model, we use a flux-averaging procedure, where the radiative transfer problem is solved for every distinct cloud configuration separately."146" By averaging these radiation fluxes weighted with the respective cloud cover values, the mean radiative flux is obtained, which then enters into the atmospheric model calculations (see Paper I for details)."," By averaging these radiation fluxes weighted with the respective cloud cover values, the mean radiative flux is obtained, which then enters into the atmospheric model calculations (see Paper I for details)."147an assumed effective temperature of 2500 K. The values given in Table 1 for the systemic velocity of the source and the terminal velocity of the outflow were based upon Herschel/HIFI observations of the CO J=6—5 transition at 691.473 GHz (Schmidt et 22010) and are in excellent agreement with previous determinations by BWO1.,an assumed effective temperature of 2500 K. The values given in Table 1 for the systemic velocity of the source and the terminal velocity of the outflow were based upon /HIFI observations of the CO $J=6-5$ transition at 691.473 GHz (Schmidt et 2010) and are in excellent agreement with previous determinations by BW01.148" To obtain an estimate of the total mass-loss rate in the outflow, we have modeled the fluxes observed by Herschel/HIFI for the CO J= 6-5, J=10-9 and J=16—15 transitions."," To obtain an estimate of the total mass-loss rate in the outflow, we have modeled the fluxes observed by /HIFI for the CO $J=6-5$ , $J=10-9$ and $J=16-15$ transitions."149 Full details of the observations and modeling of CO will be given by Schmidt et ((2010)., Full details of the observations and modeling of CO will be given by Schmidt et (2010).150" Our best estimate of the gas mass-loss rate in the inner envelope is 4.6x10°°Moyr, a factor ~3 larger than that obtained by Schóiier Olofsson (2000) from a fit to the CO J=2—I transition, and the derived CO/H> ratio is 10-3."," Our best estimate of the gas mass-loss rate in the inner envelope is $4.6 \times 10^{-6} M_\odot\, \rm yr^{-1},$ a factor $\sim 3$ larger than that obtained by Schöiier Olofsson (2000) from a fit to the CO $J=2-1$ transition, and the derived $_2$ ratio is $^{-3}$."151" Indeed, we find that a constant mass- rate model that fits the CO J=6—5, J=10—9 and J=16—15 transitions substantially overpredicts the flux in the CO J=2-1 transition."," Indeed, we find that a constant mass-loss rate model that fits the CO $J=6-5$, $J=10-9$ and $J=16-15$ transitions substantially overpredicts the flux in the CO $J=2-1$ transition."152" This discrepancy suggests some variability in the mass-loss rate, with a larger value applying to the inner envelope where the higher-lying transitions of CO originate."," This discrepancy suggests some variability in the mass-loss rate, with a larger value applying to the inner envelope where the higher-lying transitions of CO originate."153" We obtained a satisfactory fit to both the CO rotational line fluxes and the continuum spectrum by assuming a gas and dust density that decreases as , instead of the radius? density profile expected for an envelope with a constant mass loss rate and outflow velocity."," We obtained a satisfactory fit to both the CO rotational line fluxes and the continuum spectrum by assuming a gas and dust density that decreases as $^{-2.15}$, instead of the $^{-2}$ density profile expected for an envelope with a constant mass loss rate and outflow velocity."154 The gas-to-dust mass ratio in this model is 510., The gas-to-dust mass ratio in this model is 510.155" Fortunately, given the significant uncertainties in many of the assumed parameters listed in Table 1, the derived water outflow rate is not strongly dependent upon any of them."," Fortunately, given the significant uncertainties in many of the assumed parameters listed in Table 1, the derived water outflow rate is not strongly dependent upon any of them."156" As discussed in GNM, unless the mass-loss rate is extremely large, the excitation of water is dominated by radiative pumping via the 64m v» band."," As discussed in GNM, unless the mass-loss rate is extremely large, the excitation of water is dominated by radiative pumping via the $\,\mu$ m $\nu_2$ band."157" Thus, for a given water outflow rate, the observed water line fluxes scale linearly with the observed 6m continuum flux."," Thus, for a given water outflow rate, the observed water line fluxes scale linearly with the observed $\,\mu$ m continuum flux."158" Since the latter is an observed (rather than a derived) quantity, our estimate of the water outflow rate is largely independent of the distance or total outflow rate assumed for the source."," Since the latter is an observed (rather than a derived) quantity, our estimate of the water outflow rate is largely independent of the distance or total outflow rate assumed for the source."159" In modeling the water line strength and profile observed toward V Cygni, we have investigated two models for the spatial distribution of the observed water."," In modeling the water line strength and profile observed toward V Cygni, we have investigated two models for the spatial distribution of the observed water."160" In Model A, we assume that water is present at radii as small as 4.5x1014 cm, while in Model B, we adopt an inner radius Ri,=2X10cm at which the water is injected into the outflow."," In Model A, we assume that water is present at radii as small as $4.5 \times 10^{14}$ cm, while in Model B, we adopt an inner radius $R_{\rm in} = 2 \times 10^{15} \, \rm cm$ at which the water is injected into the outflow."161" In both cases, we assume an outer radius of Rout=1x10!°cm, the estimated photodissociation radius for The distribution assumed in Model B is expected if the vapourisation of icy objects is the origin ofthe observed water vapour, since all such objects at smaller distances from the star will have been vapourised already (Ford Neufeld"," In both cases, we assume an outer radius of $R_{\rm out} = 1 \times 10^{16} \, \rm cm$, the estimated photodissociation radius for The distribution assumed in Model B is expected if the vapourisation of icy objects is the origin ofthe observed water vapour, since all such objects at smaller distances from the star will have been vapourised already (Ford Neufeld"162of this diagram as an alternative to the BPT diagrams and N2 ratio to classify galaxies.,of this diagram as an alternative to the BPT diagrams and N2 ratio to classify galaxies.163 As explained in Section 3. the BPT and other optical emission lines diagnostic diagrams have become important in. the classification of galaxies.," As explained in Section 3, the BPT and other optical emission lines diagnostic diagrams have become important in the classification of galaxies."164 In this section. our aim ts to investigate the effects of the evolution of galaxies from the three BPT diagrams.," In this section, our aim is to investigate the effects of the evolution of galaxies from the three BPT diagrams."165 For this purpose. and with the objective of increase our number of galaxies. we did not take any restriction in magnitude. as detailed in the sample selection.," For this purpose, and with the objective of increase our number of galaxies, we did not take any restriction in magnitude, as detailed in the sample selection."166 In Fig., In Fig.167 5 we show the three BPT diagrams for the four redshift samples., 5 we show the three BPT diagrams for the four redshift samples.168 As redshift increases. we observe that ni] A5007/HB goes toward higher values.," As redshift increases, we observe that ] $\lambda$ goes toward higher values."169 In order to explain this shift. in Fig.," In order to explain this shift, in Fig."170 6 we plotted the ratio [Om] 415007/H8. versus redshift and metallicity only for SF galaxies selected with the Kauf03 criterion., 6 we plotted the ratio ] $\lambda$ versus redshift and metallicity only for SF galaxies selected with the Kauf03 criterion.171 The gap observed around ς~ 0.145 (see Fig., The gap observed around $z\sim$ 0.145 (see Fig.172 6a. c) is due to theHf line falling nearby the 5577 sky line. because the residuals are significant and. as a consequence. measurements of around this redshift were lost.," 6a, c) is due to the line falling nearby the 5577 sky line, because the residuals are significant and, as a consequence, measurements of around this redshift were lost."173 As shown in Fig., As shown in Fig.174 6b. there is a clear tendencyof the ui] 25007/H6 ratio towards higher values with redshift. which 1s explained by examining the same ratio against 12+log(O/H).," 6b, there is a clear tendencyof the ] $\lambda$ ratio towards higher values with redshift, which is explained by examining the same ratio against 12+log(O/H)."175 The ratio [Om] 5007/HB has demonstrated to correlate linearly with metallicity (see. for example. Liang et al.," The ratio ] $\lambda$ has demonstrated to correlate linearly with metallicity (see, for example, Liang et al."176 2006)., 2006).177 Then. a decrement in 12+log(O/H) will result in higher values of πι A5007/H (see Fig.," Then, a decrement in 12+log(O/H) will result in higher values of ] $\lambda$ (see Fig."178 6b)., 6b).179 We observe a decrement of «0.2 dex in [Om] 45007/HB. and a decrement of «0.1 dex in 12+log(O/H) for the z redshift range with respect to the zo range.," We observe a decrement of $\sim$ 0.2 dex in ] $\lambda$, and a decrement of $\sim$ 0.1 dex in 12+log(O/H) for the $z_3$ redshift range with respect to the $z_0$ range."180 In previous papers (Lara-Lóppez et al., In previous papers (Lara-Lóppez et al.181 2009a. b) we reported a decrement in. 12+log(O/H) of «0.1 dex for the redshift range 0.3«c<0.4 comparing galaxies in the same range of luminosity at different redshift intervals.," 2009a, b) we reported a decrement in 12+log(O/H) of $\sim$ 0.1 dex for the redshift range $0.3 < z < 0.4$ comparing galaxies in the same range of luminosity at different redshift intervals."182 Since the possible bias. such as Iuminosity. mass and aperture effects of those samples were carefully studied. we demonstrated there that this decrement in metallicity is due to an intrinsic evolution of the galaxies.," Since the possible bias, such as luminosity, mass and aperture effects of those samples were carefully studied, we demonstrated there that this decrement in metallicity is due to an intrinsic evolution of the galaxies."183 Although our z; sample corresponds to luminous galaxies. if we compare galaxies with the same luminosity. taking as a reference our previous papers. the metallicity decrement will be again of ~O.1 dex. and as consequence. the effects on the BPT diagrams will be the same.," Although our $z_3$ sample corresponds to luminous galaxies, if we compare galaxies with the same luminosity, taking as a reference our previous papers, the metallicity decrement will be again of $\sim$ 0.1 dex, and as consequence, the effects on the BPT diagrams will be the same."184 Therefore. the evolution observed in the in] 45007/H lines ratio toward higher values in the three BPT diagrams. could be attributed to a metallicity evolution.," Therefore, the evolution observed in the ] $\lambda$ lines ratio toward higher values in the three BPT diagrams, could be attributed to a metallicity evolution."185 On the other hand. we analyze the ratio against redshift and metallicity (see Fig.6c. d).," On the other hand, we analyze the ratio against redshift and metallicity (see Fig.6c, d)."186 The ratio is also a metallicity index. commonly known as N2. and it has been widely studied since it 15 not severely affected by dust extinction. (see Pettini & Pagel 2004).," The ratio is also a metallicity index, commonly known as N2, and it has been widely studied since it is not severely affected by dust extinction (see Pettini $\&$ Pagel 2004)."187 Among the calibrations of the N2 index. we have for example those of Raimann et al. (," Among the calibrations of the N2 index, we have for example those of Raimann et al. ("1882000). Denicoló et al. (,"2000), Denicoló et al. ("1892002). and Liang et al. (,"2002), and Liang et al. ("1902006).,2006).191 In Fig., In Fig.192 6d we observe a clear increasing trend of metallicity following the increase of the N2 index up to 12+log(O/H) «9.0., 6d we observe a clear increasing trend of metallicity following the increase of the N2 index up to 12+log(O/H) $\sim$ 9.0.193 The galaxies with 12+log(O/H) >9 show a flattening and a slightly decrease of the N2 index with metallicity (see Fig., The galaxies with 12+log(O/H) $>$ 9 show a flattening and a slightly decrease of the N2 index with metallicity (see Fig.194 6d)., 6d).195 This trend was explained by Kewley et al. (, This trend was explained by Kewley et al. (1962002) using photoionization models as follows: when the secondary production of nitrogen dominates. at somewhat higher metallicity. the line ratio continues to increase. despite the decreasing electron temperature.,"2002) using photoionization models as follows: when the secondary production of nitrogen dominates, at somewhat higher metallicity, the line ratio continues to increase, despite the decreasing electron temperature."197 Eventually. at still higher metallicities. nitrogen becomes the dominant coolant in the nebula. and the electron temperature falls sufficiently to ensure that the nitrogen line weakens with increasing metallicity.," Eventually, at still higher metallicities, nitrogen becomes the dominant coolant in the nebula, and the electron temperature falls sufficiently to ensure that the nitrogen line weakens with increasing metallicity."198 Liang et al. (, Liang et al. (1992006). using SDSS galaxies with redshifts 0.04<20.25. observed a small decrement of the N2 index against metallicity: this turnover of the N2 index is more evident for the higher redshifts z» and zi in our sample (see Fig.,"2006), using SDSS galaxies with redshifts $0.04 < z < 0.25$, observed a small decrement of the N2 index against metallicity; this turnover of the N2 index is more evident for the higher redshifts $z_2$ and $z_3$ in our sample (see Fig."200 6d)., 6d).201 The turnover of the N2 ratio will produce in the 45007/HB vs. BPT diagram. the turnover zone around N2 ~—0.4. which is more evident in a density plot (see Fig.," The turnover of the N2 ratio will produce in the ] $\lambda$ vs. BPT diagram, the turnover zone around N2 $\sim-0.4$, which is more evident in a density plot (see Fig."202 2)., 2).203 Regarding the two left BPT diagrams of Fig., Regarding the two left BPT diagrams of Fig.204 5b and c. since they share the ratio [Om] 45007/Hf. the evolutionary effects due to a decrement in metallicity will be the same as discussed above.," 5b and c, since they share the ratio ] $\lambda$, the evolutionary effects due to a decrement in metallicity will be the same as discussed above."205 The ratio has never been used before as a metallicity indicator because it is far more sensitive to ionization than to metallicity (Liang et al., The ratio has never been used before as a metallicity indicator because it is far more sensitive to ionization than to metallicity (Liang et al.206 2006)., 2006).207 Moreover. it is double-valued with metallicity (see Fig.," Moreover, it is double–valued with metallicity (see Fig."208 +). whereas the ratio is not a metallicity indicator.," 4), whereas the ratio is not a metallicity indicator."209 Therefore. after analyzing allthe ratios involved in the three BPT diagrams. we concluded that the evolution of galaxies in the three BPT diagrams is shown through the i1] 25007/HB8 ratio.," Therefore, after analyzing allthe ratios involved in the three BPT diagrams, we concluded that the evolution of galaxies in the three BPT diagrams is shown through the ] $\lambda$ ratio."210 Since this ratio is a metallicity indicator. any decrement in metallicity will result in higher values of the ui] 25007/H6 ratio.," Since this ratio is a metallicity indicator, any decrement in metallicity will result in higher values of the ] $\lambda$ ratio."211separated by less than100.! kpe. which is significantly smaller than the scales studied herein.,"separated by less than$h^{-1}$ kpc, which is significantly smaller than the scales studied herein."212" In each case. we also constructed catalogues of random data points (containing SLO"" points) over the same area of the sky and with the same selection function as discussed in Popeetal.(2004)."," In each case, we also constructed catalogues of random data points (containing $8\times10^5$ points) over the same area of the sky and with the same selection function as discussed in \cite{Pope2004}."213. These random catalogues are then used to calculate edge effects on the N-point correlation functions using the estimators presented in Szapudi&Szalay(1998)., These random catalogues are then used to calculate edge effects on the N–point correlation functions using the estimators presented in \cite{SS1998}.214 There are two common parametrizations of ().., There are two common parametrizations of $Q_z$ .215" One detines where sj». s»; and sy, are the three sides of a triangle in redshift space."," One defines where $s_{12}$, $s_{23}$ and $s_{31}$ are the three sides of a triangle in redshift space."216 Then Q(s.i.0) is defined by the ratio of the 3PCF QUS12.523.534 ). to sums of productsof 2PCFs (e.g. €(512)€(815) and permutations}:," Then $Q(s,u,v)$ is defined by the ratio of the 3PCF $\zeta(s_{12}, s_{23}, s_{31})$ , to sums of productsof 2PCFs (e.g. $\xi(s_{12})\xi(s_{13})$ and permutations):"217Wangetal.(2001a)/ outlined a method to decompose the observed polarimetry into two components.,\citet{Wang:2001} outlined a method to decompose the observed polarimetry into two components.218 On the Q - U plot. the two components correspond to the polarized vectors projected onto the so-called dominant axis and (he axis perpendicular to the dominant axis.," On the Q - U plot, the two components correspond to the polarized vectors projected onto the so-called dominant axis and the axis perpendicular to the dominant axis."219 The dominant axis can be defined rom (he aspherical distribution of the data points on the Q - U plane., The dominant axis can be defined from the aspherical distribution of the data points on the Q - U plane.220 The dominant axis is derived by a linear fit to the data points weighted bv the observational errors in the Q - U plane., The dominant axis is derived by a linear fit to the data points weighted by the observational errors in the Q - U plane.221 The spectropolarimetry projected to the dominant axis represents global geometric deviations Irom spherical svannietry. whereas the vector perpendicular to the dominant axis represents deviations from the dominant axis., The spectropolarimetry projected to the dominant axis represents global geometric deviations from spherical symmetry whereas the vector perpendicular to the dominant axis represents deviations from the dominant axis.222 The same niethod will be applied in (his study. as well., The same method will be applied in this study as well.223 We show in Figures 3 to 7 the observed data points in the Q - U plane., We show in Figures 3 to 7 the observed data points in the Q - U plane.224 Each point represent a data pair of the Q - U vector at a different wavelength., Each point represent a data pair of the Q - U vector at a different wavelength.225 The wavelength of the data points in important intervals are encoded in color., The wavelength of the data points in important intervals are encoded in color.226 The data show remarkable evolution during the five epochs of observation., The data show remarkable evolution during the five epochs of observation.227 The polarization also shows spectral features. (hal can be identified with features in the flux spectra., The polarization also shows spectral features that can be identified with features in the flux spectra.228 This firmly establishes (hat SN 2001el is intrinsically polarized. αἱ least at certain epochs alter explosion.," This firmly establishes that SN 2001el is intrinsically polarized, at least at certain epochs after explosion."229 SN 2001el exhibits some remarkable features that are unlike (hose of previously observed SN ΠΠ aad the subluminous SN Ia 1999bv (Wangetal.1996:Wang.Wheeler.&Hollich1997:Waneetal.2001a:Leonard&Filippenko2001:Howell2001 ).," SN 2001el exhibits some remarkable features that are unlike those of previously observed SN II and the subluminous SN Ia 1999by \citep{Wang:1996, WWH:1997, Wang:2001, Leonard:2001a, Leonard:2001b,230Howell:99by}."231. On the Q - U plot. SN 19955 and SN 1999by showed well-defined linear features (WangοἱILowelletal. 2001).," On the Q - U plot, SN 1998S and SN 1999by showed well-defined linear features \citep{Wang:2001,Howell:99by}."232. This is indicative of a relatively well-defined svmunetiy axes., This is indicative of a relatively well-defined symmetry axes.233 The SN 2001el data. however. show large scatter around (the dominant axis.," The SN 2001el data, however, show large scatter around the dominant axis."234 In. particular. a sharp increase of the degree of polarization is seen in the Ca 1I IK. triplet in the Sept. 26 and Oct. 1 data.," In particular, a sharp increase of the degree of polarization is seen in the Ca II IR triplet in the Sept. 26 and Oct. 1 data."235 Other strong polarized spectral features are also observed during Sept. 26 and Oct. 1., Other strong polarized spectral features are also observed during Sept. 26 and Oct. 1.236 The polarized features becomes much weaker in the data taken on Oct. 18 and afterwards., The polarized features becomes much weaker in the data taken on Oct. 18 and afterwards.237 To derive the intrinsic polarization due to the supernova atmosphere. we first need to deduce the component due to interstellar dist.," To derive the intrinsic polarization due to the supernova atmosphere, we first need to deduce the component due to interstellar dust."238 A simple approach is to assume that the resonance-scattered photons are unpolarized (Trammelletal.1993)., A simple approach is to assume that the resonance-scattered photons are unpolarized \citep{Trammell:1993}.239. This method attempts to disünguish continuum and scattered photons ancl use that separation to derive interstellar extinction (Jeffrey1991:Trammelletal.1993:Hoflich1996:WangTranetal.1997:Leonard&Filippenko 2001)..," This method attempts to distinguish continuum and scattered photons and use that separation to derive interstellar extinction \citep{Jeffrey:1991,Trammell:1993,Hoeflich:1996w,240Wang:1996, Tran:1997, Leonard:2001a}. ."241 This technique implicitly assumes a unique intrinsic, This technique implicitly assumes a unique intrinsic242pulse-heights of greater than 3.3 photoelectrons. and Cherenkov photon arrival (mes within +50 nns of the median arrival time.,"pulse-heights of greater than 3.3 photoelectrons, and Cherenkov photon arrival times within $\pm$ ns of the median arrival time."243 Clusters of at least five adjacent triggered pixels (rather than the usual four-pixel eut) were required in each event (to minimize the effects of the bright star field., Clusters of at least five adjacent triggered pixels (rather than the usual four-pixel cut) were required in each event to minimize the effects of the bright star field.244 After these pre-selection cuts. which reduced events due to background light were reduced by99%... the shower rate was stable on a run-to-run basis for observations in (he same vear.," After these pre-selection cuts, which reduced events due to background light were reduced by, the shower rate was stable on a run-to-run basis for observations in the same year."245 The systematic difference of the run-by-run acceptance within (he same vear is expected to be less than124., The systematic difference of the run-by-run acceptance within the same year is expected to be less than.246.. The ON/OFF shower rate differences in 2002 and 2003 were —8415% and —1+134... respectively.," The ON/OFF shower rate differences in 2002 and 2003 were $-8\pm15$ and $-1\pm13$, respectively."247 By examining the event rates within each run we were able to reject periods affected by cloud. dew. forming on the mirrors. instrumental abnormalities. etc.," By examining the event rates within each run we were able to reject periods affected by cloud, dew forming on the mirrors, instrumental abnormalities, etc."248 Only data taken at elevation angles greater than 607 were accepted., Only data taken at elevation angles greater than $^{\circ}$ were accepted.249 After these cuts. 4300 oof ON- and 3900 oof OFF-source data survived.," After these cuts, 4300 of ON- and 3900 of OFF-source data survived."250 ‘Trigger rates for each pixel per 700ji were monitored by a scaler circuit in real-time and recorded. each second., Trigger rates for each pixel per $\mu$ s were monitored by a scaler circuit in real-time and recorded each second.251 These data were used to exclude ‘hot’ pixels (generally due to the passage of a star through the FOV of a pixel) in off-Iime analvsis., These data were used to exclude `hot' pixels (generally due to the passage of a star through the FOV of a pixel) in off-line analysis.252 Hillas parameters were then calculated to discriminate gamnma-ravs from cosmic ravs based on the image shape aud orientation (Ilillas1985)., Hillas parameters were then calculated to discriminate gamma-rays from cosmic rays based on the image shape and orientation \citep{hillas85}.253. Further. in each νους data. we masked a small number of pixels which showed celormed ADC spectra. possibly due to a hardware fault.," Further, in each year's data, we masked a small number of pixels which showed deformed ADC spectra, possibly due to a hardware fault."254 Discrimination of the cosmic rav background from gamima-ravs was carried out using the likelihood method of Enomotoetal.(2002a)., Discrimination of the cosmic ray background from gamma-rays was carried out using the likelihood method of \cite{enomoto02a}.255. The resulting distributions of the image orientation angle. a. for the combined data in 2002 and 2003 are shown in Fig. 5..," The resulting distributions of the image orientation angle, $\alpha$, for the combined data in 2002 and 2003 are shown in Fig. \ref{fig1}."256 The normalizations between the ON- ancl were carried out using data with a>27°., The normalizations between the ON- and OFF-distributions were carried out using data with $\alpha>27^\circ$.257 The numbers of excess events 187) were 530)-120 (in an observation time of 2100 min.).," The numbers of excess events $\alpha <18^\circ$ ) were $530\pm120$ (in an observation time of 2100 min.),"258 540+140 (2200 min.).," $540\pm140$ (2200 min.),"259 and 1080X180 (4300 min.), and $1080\pm180$ (4300 min.)260" in 2002. 2003. and the combined data. respectively,"," in 2002, 2003, and the combined data, respectively."261 The excess rates for 2002 and 2003 were similar to each other., The excess rates for 2002 and 2003 were similar to each other.262 Nightly signal rates were also checked during both vears., Nightly signal rates were also checked during both years.263" The largest deviations occurred with a rate 2.50.9 times larger than the average for 2002 and 3.0£1.0 times [or 2003. respectively,"," The largest deviations occurred with a rate $2.5\pm 0.9$ times larger than the average for 2002 and $3.0\pm 1.0$ times for 2003, respectively."264 These are not unexpected statistically. and therefore there is no evidence of time-variability in the TeV emission.," These are not unexpected statistically, and therefore there is no evidence of time-variability in the TeV emission."265" To check on the spatial distribution of the signal we derived (he ""significance map. shown by the blue contours in Fig. 4.."," To check on the spatial distribution of the signal we derived the “significance map”, shown by the blue contours in Fig. \ref{fig2}."266 The contours were calculated [rom the distribution of the detection significance determined at each location [rom the difference in the a plots minus OFF-source histogram) divided by the statistical errors., The contours were calculated from the distribution of the detection significance determined at each location from the difference in the $\alpha$ plots (ON- minus OFF-source histogram) divided by the statistical errors.267 The centroid is consistent with position of the X-ray maxinmnmn. within our possible svstematic uncertainty of 0.17.," The centroid is consistent with position of the X-ray maximum, within our possible systematic uncertainty of $^\circ$ ."268,.269zIH/Ecolh(izH/E). Wi >>E. there will be a significant increase in the pair production rale.," If $H>>E$, there will be a significant increase in the pair production rate."270" The electric field of the electrosphere could be as high as E—40E,.,;;zz120 MeV?.", The electric field of the electrosphere could be as high as $E=40E_{crit}\approx 120$ $^2$.271 On the other hand (he estimated magnetic fields at the surface of the quark stars could be οἱ the order of 1210— 10G x20—2000 MeV? (1G =1.953x10.!! MeV?).," On the other hand the estimated magnetic fields at the surface of the quark stars could be of the order of $H\approx27210^{15}-10^{17}$ G $\approx 20-2000$ $^2$ (1G $=1.953\times27310^{-14}$ $^2$ )."274 Magnetic fields with such high values may be present in very voung quark stars., Magnetic fields with such high values may be present in very young quark stars.275 Assuming ecquipartition of energy. (he energy of (he differential rotation can be converted into magnetic energy. so that 0?((NO/O)zz(40/3)RCLIP/8). where 7 is the moment of inertia of the star. A its radius and Ὁ and AQ are the angular velocity and the variation of the angular velocity. respectively.," Assuming equipartition of energy, the energy of the differential rotation can be converted into magnetic energy, so that $I\Omega ^{2}\left( \Delta \Omega /\Omega \right) \approx276\left( 4\pi /3\right) R^{3}\left( H^{2}/8\pi \right) $, where $I$ is the moment of inertia of the star, $R$ its radius and $\Omega $ and $\Delta \Omega $ are the angular velocity and the variation of the angular velocity, respectively."277" Therefore the magnetic field of a voung quark star can be approximated as LIzz10!(NQ/O)!> MeV?,"," Therefore the magnetic field of a young quark star can be approximated as $H\approx27810^{4}\left( \Delta \Omega /\Omega \right) ^{1/2}$ $^{2}$."279 By assuming that AQ/Ozz0.03. we can obtain values of the magnetic field as high as 44zz2000 MeV?.," By assuming that $\Delta \Omega /\Omega \approx 0.03$, we can obtain values of the magnetic field as high as $H\approx 2000$ $^{2}$."280 Of course magnetic fields of such strength are nol stable. because thev will be pushed to and through the surface by buovant forces ancl then reconnect (IxInzniakandRuclerman1998).," Of course magnetic fields of such strength are not stable, because they will be pushed to and through the surface by buoyant forces and then reconnect \citep{Klu98}."281. For à magnetic field of the order of LFzz2000 MeV?E we have ὃμ-22.53., For a magnetic field of the order of $H\approx 2000$ $^{2}$ we have $\delta _H\approx 53$.282 Therefore strong magnetic fields can significantly increase the electron-positron pair production rate. and. consequently the Iuminositv of the electrosphere of quark stars.," Therefore strong magnetic fields can significantly increase the electron-positron pair production rate, and, consequently the luminosity of the electrosphere of quark stars."283 In the present paper we have re-considered the electron-positron pair enission from the electrosphere of quark stars. as originally proposed by Usov(1998a.b).. bv pointing out the important role (he boundary effects and (he inhomogeneity in the distribution of the electric field may play in the pair creation process.," In the present paper we have re-considered the electron-positron pair emission from the electrosphere of quark stars, as originally proposed by \citet{Us98a,Us98b}, by pointing out the important role the boundary effects and the inhomogeneity in the distribution of the electric field may play in the pair creation process."284 At zero temperature. there are no available free enerev states in (he electron. plasma at the strange stars surface.," At zero temperature, there are no available free energy states in the electron plasma at the strange star's surface."285" Therefore. at low temperatures 7'<T,220.1 MeV (corresponding to a quark star surface electric potential of V,—5 MeV). the pair production mechanism by the strong electric field of the electrosphere is severely limited bv the quantum effects aud (he exclusion principle specific to the statistics."," Therefore, at low temperatures $T\leq T_{cr}\approx 0.1$ MeV (corresponding to a quark star surface electric potential of $V_q=5$ MeV), the pair production mechanism by the strong electric field of the electrosphere is severely limited by the quantum effects and the exclusion principle specific to the Fermi-Dirac statistics."286 At high temperatures 77/77;220.1 MeV. the pair creation rate is controlled bv the electric field £ and not by the temperature. because such a process is essentially a «quantum process.," At high temperatures $T\geq T_{cr}\approx 0.1$ MeV, the pair creation rate is controlled by the electric field $E$ and not by the temperature, because such a process is essentially a quantum process."287 Once the munber of available electron states becomes higher than the schwinger pair production rate. electron-positron pairs can be freely created by (he electric field at Che surface of strange stars.," Once the number of available electron states becomes higher than the Schwinger pair production rate, electron-positron pairs can be freely created by the electric field at the surface of strange stars."288 This happens at a critical temperature {νο which strongly depends on the electrostatic properties of the quark star surface.," This happens at a critical temperature $T_{cr}$, which strongly depends on the electrostatic properties of the quark star surface."289 The critical temperature increases with the increase of the electrostatic potential Ἐν., The critical temperature increases with the increase of the electrostatic potential $V_q$.290 At high enough temperatures. the pair creation process is almost independent of the temperature and is controlled exclusively bv the electric field.," At high enough temperatures, the pair creation process is almost independent of the temperature and is controlled exclusively by the electric field."291 On the other hand. (he actual thermalized pair creation rate. which.," On the other hand, the actual thermalized pair creation rate, which,"292WOOL MNis an exteuded (~. 17) radio. source located at the tangential point (/2 197) of the Sagittarius ari.,W51 is an extended $\sim 1^\circ$ ) radio source located at the tangential point $l=49^\circ$ ) of the Sagittarius arm.293 It is composed of two complex ID II regions. W51A and W51D. and the supernova remaut (SNR) ας (es. Bicging 1975: Koo 1997).," It is composed of two complex H II regions, W51A and W51B, and the supernova remant (SNR) W51C (e.g., Bieging 1975; Koo 1997)."294 W51A forms the northern part of WL. and is separated frou the other two sources.," W51A forms the northern part of W51, and is separated from the other two sources."295 It coutains two major coniponeuts. cach composed of several compact IL II reisons.," It contains two major components, each composed of several compact H II reigons."296 W51D ids composed of at least ‘ II : scattered m area of six~DS commpacsize., W51B is composed of at least six compact H II regions scattered over an area of $\sim 15'$ size.297 W5LBHu regionssources are associatedover with a stream of atomic and molecular gases. with lue-of-sieght velocitics significauth ercater than the maxinuun velocity peruütted by Galactic rotation alone.," W51B sources are associated with a stream of atomic and molecular gases, with line-of-sight velocities significantly greater than the maximum velocity permitted by Galactic rotation alone."298 The stream is thought to be eas flowine along the Sagittarius spiral arm in response fo the perturbation the spiral poteutial (e.g... Burton 1971).," The stream is thought to be gas flowing along the Sagittarius spiral arm in response to the perturbation due to the spiral potential (e.g., Burton 1971)."299 dueSuperimposedfo ou W51D.- there is⋅ an extended structure ll.W51C which. is à SNB.," Superimposed on W51B, there is an extended structure W51C which is a SNR."300 W51C€ appears iu radio contiuuuun as an incomplete shell of ~30 extent with its upper portion open (Copetti&Schinidt1991:Subrahmanuvan&Coss 1995).," W51C appears in radio continuum as an incomplete shell of $\sim 30'$ extent with its upper portion open \citep{cop91, sub95}."301. Shocked atomic and molecular eases have Όσοι detected iu the western part of the SNR. which indicates that the SNR is interacting with a laree molecular cloud (hoo&Teiles1991:ΊνουMoon19972.b).," Shocked atomic and molecular gases have been detected in the western part of the SNR, which indicates that the SNR is interacting with a large molecular cloud \citep{koo91, koo97a, koo97b}."302. The complex. structure of W51 is partly due tfo. the inclination of the Sagittarius arii. so that we look down the leugth of the axi a distance of 5 kpc along the linc-of-sighit.," The complex structure of W51 is partly due to the inclination of the Sagittarius arm, so that we look down the length of the arm a distance of 5 kpc along the line-of-sight."303 Soft 52 keV) N-ray emission associated with N51 has been detected by aud ROSAT (Seward1990:Ίουetal.1995.hereafter.IIÍSS).," Soft $\simlt 2$ keV) X-ray emission associated with W51 has been detected by and ROSAT \citep[hereafter KKS]{sew90, koo95}."304. Diffuse N-ravs come from a region sumoundius W5ID aud W51C. whereas W51AÀ has essentially no soft N-rav eunission associated with it.," Diffuse X-rays come from a region surrounding W51B and W51C, whereas W51A has essentially no soft X-ray emission associated with it."305 The soft N-ray cutting region is elongated (50« 38) along the east-west direction. aud mav be divided into three parts: a ceutral structure composed of two," The soft X-ray emitting region is elongated $50'\times 38'$ ) along the east-west direction, and may be divided into three parts; a central structure composed of two"306axis ratio 22.,axis ratio $\simgeq \ 2$.307 Such a disk would eet quickly wound up due to differeutial rotation., Such a disk would get quickly wound up due to differential rotation.308 For these reasous we regard it unlikely that GBSN's velocity field is due to pure rotation., For these reasons we regard it unlikely that GR8's velocity field is due to pure rotation.309 Aanore likely inodoel is one iu wuch the kinematics of CRS can be described as a combination of radial aud circular notions., A more likely model is one in which the kinematics of GR8 can be described as a combination of radial and circular motions.310 Such a model provides a reasonable fit to the observed velocity field., Such a model provides a reasonable fit to the observed velocity field.311 In this interpretation. iu case the radial motions are outwards. then they could be dviven by he star formation in GR&: a previous study (Elineereenu IIunter 2000) las shown that the pressure in the TI regions in this galaxy is af least 55 times greater than he average pressure in the disk.," In this interpretation, in case the radial motions are outwards, then they could be driven by the star formation in GR8; a previous study (Elmegreen Hunter 2000) has shown that the pressure in the HII regions in this galaxy is at least $55$ times greater than the average pressure in the disk."312 The measured expansion velocity is considerably less than the estimated escape velocity. so even in this iuterpretation the cold eas is still vound to the ealaxy.," The measured expansion velocity is considerably less than the estimated escape velocity, so even in this interpretation the cold gas is still bound to the galaxy."313 Finally. the radial motions could also ve immterpreted as iufall. in which case CRS is either iu the xocess of formation. or the ISM is falling |ck after a ouevious pase of expansion.," Finally, the radial motions could also be interpreted as infall, in which case GR8 is either in the process of formation, or the ISM is falling back after a previous phase of expansion."314half maximum of 0.5ddex. in the stellar. radii around an isochronal value.,"half maximum of dex, in the stellar radii around an isochronal value."315" ""This result is almost. independent of which. evolutionary models are. considered.", This result is almost independent of which evolutionary models are considered.316 An obvious interpretation is that this is caused by a spread in ages and this is investigated in the next subsection., An obvious interpretation is that this is caused by a spread in ages and this is investigated in the next subsection.317 Two separate models for an age spread. were tested., Two separate models for an age spread were tested.318" The first was a Gaussian distribution of log), age around a central. isochrone.", The first was a Gaussian distribution of $\log_{10}$ age around a central isochrone.319" The free parameters were the central age and the age dispersion (in logarithmic units). 0,."," The free parameters were the central age and the age dispersion (in logarithmic units), $\sigma_a$."320 The second model was an age distribution which is zero up to some starting age. jumps to a maximum and then decays exponentially with à decay constant. A. expressed in Myr.," The second model was an age distribution which is zero up to some starting age, jumps to a maximum and then decays exponentially with a decay constant, $\lambda_a$, expressed in Myr."321 This latter model. with a suitably small starting age. represents the exponentially accelerating star formation model advocated by Palla Stahler (1999).," This latter model, with a suitably small starting age, represents the exponentially accelerating star formation model advocated by Palla Stahler (1999)."322 The lowest starting age considered was 0.03 δλδνε., The lowest starting age considered was $0.03$ Myr.323 Ages drawn randomly from the Gaussian. age distribution were transformed into radii using the appropriate stellar models (at. the Zi of each star in the observational dataset) ancl these radii were perturbed according {ο the I50KK Zr uncertainties and then subjected to the measurement. uncertainies. random axial orientations and. selection ellects before comparing the observed and modelled clistribution of 2sinz/Itasis.," Ages drawn randomly from the Gaussian age distribution were transformed into radii using the appropriate stellar models (at the $T_{\rm eff}$ of each star in the observational dataset) and these radii were perturbed according to the K $T_{\rm eff}$ uncertainties and then subjected to the measurement uncertainties, random axial orientations and selection effects before comparing the observed and modelled distribution of $R\sin i/R_{{\rm 3Myr}}$."324 A grid of models covering a wide range of cenral ages and age dispersion was calculated. for both. the DAMOT7 ancl SOO models., A grid of models covering a wide range of central ages and age dispersion was calculated for both the DAM97 and S00 models.325 Assuming that the correct solution lay within this erid. | normalised the Ίντο probabilities ancl this gave a pair of relative probability. grids. which are shown in," Assuming that the “correct” solution lay within this grid, I normalised the K-S probabilities and this gave a pair of relative probability grids which are shown in"326no simultaneous CCD 2 observations had. been obtained. were utilized.,"no simultaneous CCD $B$ observations had been obtained, were utilized."327 Vhe datasets of the variables used. for. the »eriod. change study are listed in Table 4.., The datasets of the variables used for the period change study are listed in Table \ref{felhasznalt}.328 For each star. the dataset that defines the mean (master) light curve is marked ον an asterisk.," For each star, the dataset that defines the mean (master) light curve is marked by an asterisk."329 Phe complete table is available as supporting information in the electronic version of this article., The complete table is available as supporting information in the electronic version of this article.330 The mean light curves were constructed in the form of ifth and seventh order Fourier series for the Rite and RRab stars. respectively.," The mean light curves were constructed in the form of fifth and seventh order Fourier series for the RRc and RRab stars, respectively."331 In case of noisy and/or strongly variable ight curves. lower order fits were applied to avoid. unreal. wavy mean light-curve shapes.," In case of noisy and/or strongly variable light curves, lower order fits were applied to avoid unreal, wavy mean light-curve shapes."332 As the first step. cach dataset was. magnitude iomogenized by fitting the mean light curve.," As the first step, each dataset was magnitude homogenized by fitting the mean light curve."333 In the second step. the mean light curve was fitted in phase to 24 vears ong segments of the magnitude homogenized. combined ight curve.," In the second step, the mean light curve was fitted in phase to 2–4 years long segments of the magnitude homogenized, combined light curve."334 By default. the data were divided. into 19 segments. and their phase-shift. values corresponded to the OC for the mean epoch of the data subsets.," By default, the data were divided into 19 segments, and their phase-shift values corresponded to the $O-C$ for the mean epoch of the data subsets."335 Some of the xns were divided into shorter segments in case of very fast »eriod. changes. or they were merged together if insullicient number of cata points were available in one of the bins. and no significant period change occurred in the given time interval.," Some of the bins were divided into shorter segments in case of very fast period changes, or they were merged together if insufficient number of data points were available in one of the bins, and no significant period change occurred in the given time interval."336 The O—C values. derived in this way. were used to construct the phase-shift diagrams (O6 diagrams).," The $O-C$ values, derived in this way, were used to construct the phase-shift diagrams $O-C$ diagrams)."337 Since there were gaps in the coverage of the observations. evele count ambiguities occurred. in some cases. especially for stars showing irregular. Large period. changes.," Since there were gaps in the coverage of the observations, cycle count ambiguities occurred in some cases, especially for stars showing irregular, large period changes."338 These οC points were adjusted. upward or downward. by integer evele numbers to obtain the most acceptable period-change solution (see details in the next section)., These $O-C$ points were adjusted upward or downward by integer cycle numbers to obtain the most acceptable period-change solution (see details in the next section).339 Then. he magnitude homogenization was refined. as described in section 2.3 of Szeidletal.(2011).," Then, the magnitude homogenization was refined as described in section 2.3 of \citet{m5}."340. For the stars with abrupt »eriod changes. this step greatly improved. the magnitude 10mogenization.," For the stars with abrupt period changes, this step greatly improved the magnitude homogenization."341 The determination of the O6 points and he construction of the O—C' diagrams were finalized on the refined. version of the magnitude-homogenized data., The determination of the $O-C$ points and the construction of the $O-C$ diagrams were finalized on the refined version of the magnitude-homogenized data.342 The O6 diagrams were fitted bv polynomials. and the xeriod-change rates of the variables were measured by the cocllicients of these functions.," The $O-C$ diagrams were fitted by polynomials, and the period-change rates of the variables were measured by the coefficients of these functions."343 The reliability of the OC its were checked by the comparison of the results of direct »eriod determinations with the derivative of the O6 fits. and by the construction of folded light curves corrected. for he period change.," The reliability of the $O-C$ fits were checked by the comparison of the results of direct period determinations with the derivative of the $O-C$ fits, and by the construction of folded light curves corrected for the period change."344 Table 5. gives an example of the results for VI., Table \ref{ocvalues} gives an example of the results for V1.345 The ime interval. the corresponding mean epoch. the number of datapoints of the subsets and the O or temporal »eriod. values and their errors are given in the columns.," The time interval, the corresponding mean epoch, the number of datapoints of the subsets and the $O-C$ or temporal period values and their errors are given in the columns."346 The complete table. including the O6) values and instantaneous »eriods for all the analvzed stars is available in the electronic version of the article as supporting information.," The complete table, including the $O-C$ values and instantaneous periods for all the analyzed stars is available in the electronic version of the article as supporting information."347 The results of the period-change analysis are documentec in Fig., The results of the period-change analysis are documented in Fig.348 1. for 129 Ht Lyrae stars., \ref{oc} for 129 RR Lyrae stars.349 The variables are shown in the order of the length of their periods., The variables are shown in the order of the length of their periods.350 Decause of their extremely strong phase modulations (Blazhko effect). ane abrupt changes of the pulsation period. no period-change solution could be obtained. for three RRab variables (V5. V50 and V130).," Because of their extremely strong phase modulations (Blazhko effect) and abrupt changes of the pulsation period, no period-change solution could be obtained for three RRab variables (V5, V50 and V130)."351 The strong changes of their modal conten made the analysis of V79 and. VOO impossible using our method., The strong changes of their modal content made the analysis of V79 and V99 impossible using our method.352 Thus. οC' diagrams have not been constructec for five of the studied: 134 variables.," Thus, $O-C$ diagrams have not been constructed for five of the studied 134 variables."353 Three panels are shown for each variable., Three panels are shown for each variable.354 The lelt-harn panels show the O6 diagrams ancl their polynomial fits., The left-hand panels show the $O-C$ diagrams and their polynomial fits.355" Variables showing Blazhko effect or double-mocde pulsations are denoted as ""DE. and cdm. respectively."," Variables showing Blazhko effect or double-mode pulsations are denoted as `Bl' and `dm', respectively."356 The dillerences. between temporal period values (P?) etermined by direct period analysis of the data subsets and 10 mean period (41) adopted to caleulate O—C' values are Xotted in the middle panels (2P.)107 d]., The differences between temporal period values $P$ ) determined by direct period analysis of the data subsets and the mean period $P_a$ ) adopted to calculate $O-C$ values are plotted in the middle panels $(P-P_a)\times 10^5$ d].357 Phe variations X the instantaneous periods {0} predicted from the OC ata using Eq., The variations of the instantaneous periods $P(t)$ ] predicted from the $O-C$ data using Eq.358 3 of Juresiketal.(2001) are also shown or comparison., 3 of \citet{ocen} are also shown for comparison.359 When evele-count ambiguity of the ο6 Pvalues occurred. that solution was accepted. which vielded 1e smallest. cdillerence of the directly observed. period values rom the 2) function.," When cycle-count ambiguity of the $O-C$ values occurred, that solution was accepted, which yielded the smallest difference of the directly observed period values from the $P(t)$ function."360 The errorbars shown for the ο6 points and the observed: period values indicate. 20 formal uncertainties etermined. by least-squares fitting method., The errorbars shown for the $O-C$ points and the observed period values indicate $2\sigma$ formal uncertainties determined by least-squares fitting method.361 The right-hand panels show the folded. light curves of 1e time-transformec (Iq., The right-hand panels show the folded light curves of the time-transformed (Eq.362 4 of Juresiketal. (2001))) data. which are corrected for the period. variations according to 16 polynomial fits of the O—C' diagrams.," 4 of \citet{ocen}) ) data, which are corrected for the period variations according to the polynomial fits of the $O-C$ diagrams."363 The OC solutions are shown in two separate plots using different periods for V13 anc ΝΕ. as the period anges of these stars are too large to be shown in a single plot.," The $O-C$ solutions are shown in two separate plots using different periods for V13 and V41, as the period changes of these stars are too large to be shown in a single plot."364 The very irregular O—C' variations of thirteen stars (V12. V17T. VIS. Ves. V34. V35. V43. V44. V4T. V54. V1O. Vso and ΧΙΟ) could be fitted. only by two consecutive. separate polvnomials.," The very irregular $O-C$ variations of thirteen stars (V12, V17, V18, V28, V34, V35, V43, V44, V47, V54, V70, V80 and V110) could be fitted only by two consecutive, separate polynomials."365 The period-change and light-curve characteristics of the analyzed stars are summarized in Table 6.., The period-change and light-curve characteristics of the analyzed stars are summarized in Table \ref{tabla}.366 The first three columns list the designation. the tvpe of the variable. anc the period. used for the construction of the O6 diagram.," The first three columns list the designation, the type of the variable, and the period used for the construction of the $O-C$ diagram."367 The next two columns give photometric information on the light curves: Vi is the intensityv-averaged mean 1. magnitude. and zd is the full amplitude in the V band.," The next two columns give photometric information on the light curves: $V_\mathrm{i}$ is the intensity-averaged mean $V$ magnitude, and $A_V$ is the full amplitude in the $V$ band."368"are the Alaxwell-Boltzmann distribution. the Ixnudsen number. and Aj""in=ey,fy, ts. a thermal electron. collision_ mean [ree path.","are the Maxwell-Boltzmann distribution, the Knudsen number, and $\lambda_{e}^{th}=v_{th}/\nu_{th}$ is a thermal electron collision mean free path."369 Note that only /ALB is still dimensional but all other quantities- are cimiensionless., Note that only $f^{MB}$ is still dimensional but all other quantities are dimensionless.370. ⊳∖⋠⋠To eliminate £4“De! [rom eq., To eliminate $E^{DC}$ from eq.371 12 we need to relate it to the spatial variability of /MMOBO-, \ref{eq:fir_or} we need to relate it to the spatial variability of $f^{MB}$.372 cqToward that goal one uses the current free. condition., Toward that goal one uses the current free condition.373nc In astrophysical conditions. thermoelectric fields. would cause the plasma to settle down to à zero current. state.," In astrophysical conditions, thermoelectric fields would cause the plasma to settle down to a zero current state."374 Vhe current free condition reads: Using eq. 12..," The current free condition reads: Using eq. \ref{eq:fir_or},"375 in eq., in eq.376 14. and substituting the result into eq., \ref{eq:day_kvar} and substituting the result into eq.377 12. we find after some algebra that the current [ree distribution function is: and the ratio of Y2/Y* should be evaluated once the value of £ has been determined.," \ref{eq:fir_or}378 we find after some algebra that the current free distribution function is: where and the ratio of $Y^{2}/Y^{1}$ should be evaluated once the value of $\tilde{\nu}$ has been determined."379 For a collision dominated plasma the value of @ is given by eq. 2.., For a collision dominated plasma the value of $\tilde{\nu}$ is given by eq. \ref{eq:coll_coll}.380 Thus. the value of Y7/Y! can be evaluated.," Thus, the value of $Y^{2}/Y^{1}$ can be evaluated."381 However. if the plasma is whistler dominated. fis a functional of ο. and not a simple function obit.," However, if the plasma is whistler dominated, $f$ is a functional of $\tilde{\nu}$, and not a simple function of it."382 Substituting 2=ὃ into eq. 16.," Substituting $\tilde{\nu}=\tilde{v}^{-3}$ into eq. \ref{eq:cur_ints},"383 one linds 37/1!=4 and using this in eq. 15..," one finds $Y^{2}/Y^{1}=4 $ and using this in eq. \ref{eq:tcur_free},"384 one obtains MM ⋠⋠ . . T∐∐⊳∖⊳∖∪↓⋯↓∪⊔↓⊳∖∖⇁⋜↧↓⊓⇂↓≻↓⋅∪∖⇁⊔⇂⋖⊾∠⇂↿⇂⋯↥↙∣∶↙ ⇉↓∪∶↿∖∠⇂⋯⋅↥⇂↥∢⋅∖⇁∢⊾↓⋯⋰↓↿∙∖⇁ dependence of the collisions mean free path: Grav Willkenny 1980).," one obtains This solution is valid provided that $\epsilon^{th} =\epsilon^{coll}(v^{th})385\le 2 \times 10^{-2}$ (due the velocity dependence of the collisions mean free path; Gray Killkenny 1980)."386" Note that if (L210? then [or =2c""0.32ας1. for P=4 the uncerlving assumptions of the expansion do not hold."," Note that if $\epsilon^{th}=2\times10^{-2}$ then for $\tilde{v}=2\Rightarrow\epsilon^{coll}=0.32\not\ll 1$, for $\tilde{v}=4$ the underlying assumptions of the expansion do not hold."387 The aim of this section is to a) calculate when eq., The aim of this section is to a) calculate when eq.388 17 gives rise to whistlers. b) to find the angular dependence of the whistler spectrum. and ο) with this whistler spectrum to estimate the mocified steady state cistribution function (to which the collision dominated plasma distribution function evolves). from which follows the heat [lux inhibition factor.," \ref{eq:knud} gives rise to whistlers, b) to find the angular dependence of the whistler spectrum, and c) with this whistler spectrum to estimate the modified steady state distribution function (to which the collision dominated plasma distribution function evolves), from which follows the heat flux inhibition factor."389 Consistency checks of the assumptions made are carried out as we progress., Consistency checks of the assumptions made are carried out as we progress.390 “Lo obtain these goals we consider approximate solutions to the equations given in 7? under the assumption that whistlers pervade the plasma.," To obtain these goals we consider approximate solutions to the equations given in \ref{sec:gov}391 under the assumption that whistlers pervade the plasma."392 However. before we do so we need to establish when this situation would occur and what the qualitative nature of it is.," However, before we do so we need to establish when this situation would occur and what the qualitative nature of it is."393 The formal expression for the distribution function has been found via the Knudsen expansion and is given by eq. 15.., The formal expression for the distribution function has been found via the Knudsen expansion and is given by eq. \ref{eq:tcur_free}.394 To calculate when whistlers are present in the plasma we need to use eq., To calculate when whistlers are present in the plasma we need to use eq.395 15 in eq. LO., \ref{eq:tcur_free} in eq. \ref{eq:wisgro0}.396 Ht proves mathematically convenient to formulate all expressions in terms of Ay=hy and X as the independent variables rather than & ancl y., It proves mathematically convenient to formulate all expressions in terms of $k_{\parallel}=k\chi$ and $\chi$ as the independent variables rather than $k$ and $\chi$.397 After some algebra we obtain the following schematic form for the quasilincar growth rate: ⋜⋯∠⇂∫↗↴⊲↓⊳∖⇂↓↥∢⊾∢⊾↓⋖⋅≼⇍↿↓⋅∪⊔⊳∖↓≻↓⋅∢⊾⊳∖⊳∖⊔↓⋅⋖⋅⊳⋜⋃⊔⇂∫↗⇂↓↕∢⋅⋯⋯↓⋏∙≟∥⊳∖ pressure.," After some algebra we obtain the following schematic form for the quasilinear growth rate: where and $P_{e}$ is the electrons pressure, and $P$ the total gas pressure."398 We have chosen the magnetic field. pointing from hot to cold., We have chosen the magnetic field pointing from hot to cold.399 For this choice whistler propagation is parallel to field lines. in the coldward direction.," For this choice whistler propagation is parallel to field lines, in the coldward direction."400 Thus. for wave growth we must have y20 (or hy> 0).," Thus, for wave growth we must have $\chi>0$ (or $\tilde{k}_{\parallel}>0$ )."401 This point is taken into account in the expression written above., This point is taken into account in the expression written above.402 Lhe terms QUmU aud PUv) in eq.," The terms $Q(\tilde{k}_{\parallel},\chi)$ and $P(\tilde{k}_{\parallel},\chi)$ in eq."403 19. represent a growth term due to the OlOe term. and the damping term due to the erpορ term in eq. LO.," \ref{eq:P_Q}404 represent a growth term due to the ${\partial {F} \over \partial \mu}$ term, and the damping term due to the ${\partial {F} \over \partial405p}$ term in eq. \ref{eq:wisgro0},"406 respectively., respectively.407 llore we show that under a wide range of conditions. a distribution function. were it. &overned. only by collisions. would be unstable to whistler waves which. of course. mocitv the distribution function.," Here we show that under a wide range of conditions, a distribution function, were it governed only by collisions, would be unstable to whistler waves which, of course, modify the distribution function."408 Using eq., Using eq.409 18 we first establish when a collision dominated. plasma becomes unstable., 18 we first establish when a collision dominated plasma becomes unstable.410" Using 1Y7/Y*=+ ancl p=IEPETI""= ὰ eg. 19."," Using $Y^{2}/Y^{1}=4 $ and $\tilde{\nu}^{w,s}=\tilde{\nu}^{w,a}=0$ in eq. \ref{eq:P_Q},,"411 we have:, we have:412"range 107?—10-?, and imply that a planet can open a gap at 1 AU only if its mass is larger than about 0.1 M.","range $10^{-2}-10^{-3}$, and imply that a planet can open a gap at 1 AU only if its mass is larger than about 0.1 $M_J$."413" To open a gap at 30 AU, the mass must be larger than about 0.5 Mj."," To open a gap at 30 AU, the mass must be larger than about 0.5 $M_J$."414" To investigate the effects that a planet more massive than 0.1 M; might have on the observations of the dust continuum emission, we simulated the presence of a planet in the DG Tau disk by opening a gap in the surface density distribution corresponding to the best fit models discussed above."," To investigate the effects that a planet more massive than 0.1 $M_J$ might have on the observations of the dust continuum emission, we simulated the presence of a planet in the DG Tau disk by opening a gap in the surface density distribution corresponding to the best fit models discussed above."415" For simplicity, we assumed that the planet describes a circular orbit and that the gap can be represented by a circular ring."," For simplicity, we assumed that the planet describes a circular orbit and that the gap can be represented by a circular ring."416" To be compatible with numerical simulations of planet-disk interaction, the half-width of the ring A is assumed to be equal to twice the Hill radius Ry=Rp*/Mp/(3M,) (e.g.Brydenetal.1999;Wolfetal. 2007)."," To be compatible with numerical simulations of planet-disk interaction, the half-width of the ring $\Delta$ is assumed to be equal to twice the Hill radius $R_H = R_p \sqrt[3]{M_P/(3M_{\star})}$ \citep[e.g.][]{br99,wo07}."417". In the region between R,+A the surface density is depleted by a fraction f that depends on the mass of the planet and on the disk viscosity.", In the region between $R_p \pm \Delta$ the surface density is depleted by a fraction $f$ that depends on the mass of the planet and on the disk viscosity.418" For a= 10-3, we can assume f—0 for planet masses M,>1M;, f=0.1 for M,=0.5 My, f=0.17 for M,=0.3M; and f=0.6 for M,=0.1M; (Wolfetal.2007)."," For $\alpha = 10^{-3}$ , we can assume $f=0$ for planet masses $_{p} > 1 M_J$, $f=0.1$ for $M_p=0.5$ $M_J$, $f=0.17$ for $M_p=0.3~M_J$ and $f=0.6$ for $M_p=0.1~M_J$ \citep{wo07}."419". Therefore, only planets more massive than 1 M; will produce completely cleaned gaps."," Therefore, only planets more massive than 1 $M_J$ will produce completely cleaned gaps."420 We simulated gaps corresponding to planets in the mass range 0.3-5 M; and with orbital radii between 1 and 90 AU., We simulated gaps corresponding to planets in the mass range 0.3-5 $M_J$ and with orbital radii between 1 and 90 AU.421 For each model we calculated the residuals as the difference between the observations of DG Tau at 1.3 mm and the model image., For each model we calculated the residuals as the difference between the observations of DG Tau at 1.3 mm and the model image.422 If the gap istoo small compared to, If the gap istoo small compared to423accurate measure of the PN content of the Galaxy is important in determining the relative importance of the various routes to PN formation and the role of binary interaction (see de Marco (2009) for a review).,accurate measure of the PN content of the Galaxy is important in determining the relative importance of the various routes to PN formation and the role of binary interaction (see de Marco \cite{2009PASP..121..316D} for a review).424" It has been noted that --emitting PN tend to lie at low Galactic latitude, with a scale height of ppc (Kastner et al. (1996)))"," It has been noted that -emitting PN tend to lie at low Galactic latitude, with a scale height of pc (Kastner et al. \cite{1996ApJ...462..777K}) )"425" and the same trend exists for the pre-PN eemitters, which tend to be bipolar (Kelly Hrivnak (2005)))."," and the same trend exists for the pre-PN emitters, which tend to be bipolar (Kelly Hrivnak \cite{2005ApJ...629.1040K}) )."426" A deep Galactic plane ssurvey, such as UWISH2, has the potential to uncover a significant population of evolved objects and its contiguous areal coverage will provide important constraints on their space density and any variation along the plane."," A deep Galactic plane survey, such as UWISH2, has the potential to uncover a significant population of evolved objects and its contiguous areal coverage will provide important constraints on their space density and any variation along the plane."427" The UWISH2 survey also provides a unique study of PN morphologies (note the considerable improvement in spatial resolution over Spitzer evident in refirdc)), particularly when complemented by IPHAS optical and CORNISH radio survey data."," The UWISH2 survey also provides a unique study of PN morphologies (note the considerable improvement in spatial resolution over Spitzer evident in \\ref{irdc}) ), particularly when complemented by IPHAS optical and CORNISH radio survey data."428 Point-symmetric structure is seen in many post-AGB objects and PN; this is thought to result from the interaction of a precessing jet with the remnant AGB shell (Kwok (2000)))., Point-symmetric structure is seen in many post-AGB objects and PN; this is thought to result from the interaction of a precessing jet with the remnant AGB shell (Kwok \cite{2000eaa..bookE5200K}) ).429" In many ways the physics of post-AGB evolution mimics that seen in pre-main-sequence objects, something that could be investigated further with these simultaneous observations of both classes of object."," In many ways the physics of post-AGB evolution mimics that seen in pre-main-sequence objects, something that could be investigated further with these simultaneous observations of both classes of object."430 Observations of the near-IR aand CO bandhead emission in the post-AGB object 118276-1431 show striking similarities with emission features in Herbig, Observations of the near-IR and CO bandhead emission in the post-AGB object 18276-1431 show striking similarities with emission features in Herbig431"Sincee Tig1XR7 2rfor coustan Iuniuosity. his ""Mis also a scaling to the Newtonian gravity of the models.","Since $\Teff ^4 \propto R_*^{-2}$ for constant luminosity, this is also a scaling to the Newtonian gravity of the models."432" Figure 8 shows the normalized gi, of Fe for he moclels A (top) and B (bottom).", Figure \ref{f_ABFe} shows the normalized $g_{\rm L}$ of Fe for the models A (top) and B (bottom).433 The right hand figures show au chlarecinent of the region near the sonic poit., The right hand figures show an enlargement of the region near the sonic point.434 It shows that for model D gy of Fe around the sonic point is more than a factor two larger than for nodel A (see Fies. S((, It shows that for model B $g_{\rm L}$ of Fe around the sonic point is more than a factor two larger than for model A (see Figs. \ref{f_ABFe}( (435b) aud (dj).,b) and (d)).436 This extra amount of Fe din model Bcauses an increase in the tote? gp m the subsonic part of the wind also. as cau be secu in Fie. 9((," This extra amount of Fe in model Bcauses an increase in the $total$ $g_{\rm L}$ in the subsonic part of the wind also, as can be seen in Fig. \ref{f_ABCABC}( ("437b).,b).438 Now the effect of gp ou ex will be examined., Now the effect of $g_{\rm L}$ on $\vinf$ will be examined.439" Therefore. Model D is compared to model €. We remind that mocels D aud € have the sameZig.. aud hence the same radiative flux aud gravity. but model € las a twice as small value of Vaγιος 05 model D. Figure 9((a) slows the normalized gr, for models A. Baud €. As expected. gp(re) for model € is significantly smaller than gi(7) for models A aud B. This is Obviously due to the simaller value of."," Therefore, Model B is compared to model C. We remind that models B and C have the same, and hence the same radiative flux and gravity, but model C has a twice as small value of $\ratio$ as model B. Figure \ref{f_ABCABC}( (a) shows the normalized $g_{\rm L}$ for models A, B and C. As expected, $g_{\rm L}(r)$ for model C is significantly smaller than $g_{\rm L}(r)$ for models A and B. This is obviously due to the smaller value of."440. The iuteeral Γι)de in Fig. 9((, The integral $\int g_{\rm L}(r)~dr$ in Fig. \ref{f_ABCABC}( (441"a) for model A and D is larger than for model C. The values of fgi(r)dr for the models are 2.31 1010 aud 1.92 & 1010 ene 7 for models A aud D respectively,. aud 6.12Lae « 10175 cm?2 7 2ofor model C. Using7 Eq.","a) for model A and B is larger than for model C. The values of $\int g_{\rm L}(r)~dr$ for the models are 2.34 $\times$ $^{16}$ and 1.92 $\times$ $^{16}$ $^2$ $^{-2}$ for models A and B respectively, and 6.12 $\times$ $^{15}$ $^2$ $^{-2}$ for model C. Using Eq."442 7 mx the values of ffrom cohuun (61) iu Table 1.. the output values for cean be obtained from the values of the iutegral of gr.," \ref{eq:vinfty} and the values of from column (4) in Table \ref{t:parameters}, the output values for can be obtained from the values of the integral of $g_{\rm L}$."443 The derived output values for ffor the moclels are == 2050. 1860 aud 920 kaw | respectively for the models A.D and C. These values are equal within 10 to the oeiput values for Ισ were indicated in column (5) of Table 1..," The derived output values for for the models are = 2050, 1860 and 920 km $^{-1}$ respectively for the models A,B and C. These values are equal within 10 to the input values for which were indicated in column (5) of Table \ref{t:parameters}."444 We can conclude that a smaller value for Hs indeed cousistent with a smaller value of the integral {ο(ή) dr., We can conclude that a smaller value for is indeed consistent with a smaller value of the integral $\int g_{\rm L}(r)~dr$ .445 However. this is not au independent check. since the calculated line acceleration of optically thick lines (in the Sobolev approximation) is iuversely proportional to the Sobolev optical depth which is," However, this is not an independent check, since the calculated line acceleration of optically thick lines (in the Sobolev approximation) is inversely proportional to the Sobolev optical depth which is"446has modeled the distance of the termination shock (TS) in the direction of the Vovager 1 motion. and found that 1.5 vvields a TS distance of 94 AU. which agrees with the recently measured Vovager 1 value (81)).,"has modeled the distance of the termination shock (TS) in the direction of the Voyager 1 motion, and found that $\sim$ 1.5 yields a TS distance of 94 AU, which agrees with the recently measured Voyager 1 value \ref{sec:intro}) )."447 I argued that aat (he Sun should be similar to the uniform component οἱ iinferred from pulsar data. Biy|~ 1.6G.. which dominates low density interarm regions such as surrounding the Sun1990).," I argued that at the Sun should be similar to the uniform component of inferred from pulsar data, $\sim$ 1.6, which dominates low density interarm regions such as surrounding the Sun."448. ILowever stronger fields are indicated. 2.6 [IG if equipartition between thermal and magnetic pressure applies.," However stronger fields are indicated, $\sim$ 2.6, if equipartition between thermal and magnetic pressure applies."449 The value 11.5 wwill be used in the following discussions., The value 1.5 will be used in the following discussions.450 The interaction between ISDGs and the heliosphere depends on the dust charge. mass. and composition.," The interaction between ISDGs and the heliosphere depends on the dust charge, mass, and composition."451 In this section the results of radiative transfer models of the local ISM are compared wilh a reference abundance lor the ISM. here assume to be solar abundances. to determine the dust composition.," In this section the results of radiative transfer models of the local ISM are compared with a reference abundance for the ISM, here assume to be solar abundances, to determine the dust composition."452 The following section gives the gas-to-dust mass ratio calculated from the same assumptions., The following section gives the gas-to-dust mass ratio calculated from the same assumptions.453 Unfortunately. both solar abundances aud (he ISM composition. generally assumed to be (he summed abundances of the gas and dust. ave highly uncertain.," Unfortunately, both solar abundances and the ISM composition, generally assumed to be the summed abundances of the gas and dust, are highly uncertain."454 For instance. estimates of the solar ratio for Ο/Η vary by 230542005).," For instance, estimates of the solar ratio for O/H vary by $\sim$."455.. Also. eas and grains may decouple in transient violent interstellar phenomena2004).," Also, gas and grains may decouple in transient violent interstellar phenomena."456. The predicted LIC gas-phase abundanees of C. N. ο. Ale. Al. Si. and Fe are listed in Table 1 for the best-fitting RT models 2 and 8 (SE02).," The predicted LIC gas-phase abundances of C, N, O, Mg, Al, Si, and Fe are listed in Table \ref{tab:dust} for the best-fitting RT models 2 and 8 (SF02)."457 Comparisons between the gaseous Fe. Me. Si. and O abundances and solar abundances then vield underabuncances of Fe. Meg. Si. and O in the eas within ~1 pe of the Sun towards e CMa," Comparisons between the gaseous Fe, Mg, Si, and O abundances and solar abundances then yield underabundances of Fe, Mg, Si, and O in the gas within $\sim$ 1 pc of the Sun towards $\epsilon$ CMa."458 For this discussion. (the abundances presented in are utilized.," For this discussion, the abundances presented in are utilized."459 The short length of the local ISM towards € CMa (1 pe) contains two velocity components separated bv ~8 , The short length of the local ISM towards $\epsilon$ CMa $\sim$ 1 pc) contains two velocity components separated by $\sim 8$ 460Although these secondary GeV eamuna-ravs are not taken iuto account here aud it is beyond the scope of this paper. it is naportaut to study how laree is this spectral distortion iu future work.,"Although these secondary GeV gamma-rays are not taken into account here and it is beyond the scope of this paper, it is important to study how large is this spectral distortion in future work."461 We here discuss the expected properties of gamma-ray clusters of galaxies., We here discuss the expected properties of gamma-ray clusters of galaxies.462" Perhaps the most natural question in this regard would be ""Are they already observed iu other wavebauds such as x-rays or optical surveys?”", Perhaps the most natural question in this regard would be “Are they already observed in other wavebands such as x-rays or optical surveys?”463 We have checked that there is no statistically significant association of the ROSAT Brightest Cluster Sample (RBCS. Ebeliug et al.," We have checked that there is no statistically significant association of the ROSAT Brightest Cluster Sample (RBCS, Ebeling et al."464 1998). within the error circles of the nnideutified sources with |b]230° in the EGRET catalog.," 1998), within the error circles of the unidentified sources with $|b|>30^\circ$ in the EGRET catalog."465 We have also checked the correlation with the clusters in the revised Abell catalog (Abell Corwin. Olowin 1989). and uo statistically significant associations are found. either.," We have also checked the correlation with the clusters in the revised Abell catalog (Abell, Corwin, Olowin 1989), and no statistically significant associations are found, either."466 ILlowever. iu the following[m] we argueOo that the Ooσαταν clusters proposed in this paper are very difficult to detect in xaavs or optical bands compared with ordinary clusters ideuti&ed in these wavebands. and heuce our scenario 1s uot rejected bv these results.," However, in the following we argue that the gamma-ray clusters proposed in this paper are very difficult to detect in x-rays or optical bands compared with ordinary clusters identified in these wavebands, and hence our scenario is not rejected by these results."467 We first estimate the expected. xay flux from gamuna-rav clusters., We first estimate the expected x-ray flux from gamma-ray clusters.468 Barvouic gas in most clusters of galaxies observed in X-ravs seenmis to be in approximate livdrostatic equilibrimu with the surface brightuess well fitted by a density profile. pear)x|ld(en|Ἐ (ee. Sarazin 1988). where r. is the core radius that is typically about ~10 times smaller than the virial radius.," Baryonic gas in most clusters of galaxies observed in x-rays seems to be in approximate hydrostatic equilibrium with the surface brightness well fitted by a density profile, $\rho_{\rm gas}(r) \propto469[1+(r/r_c)^2]^{-1}$ (e.g., Sarazin 1988), where $r_c$ is the core radius that is typically about $\sim 10$ times smaller than the virial radius."470 Since the x-ray cuussivity ds proportional to D the x-ray cussion is strongly concentrated into the central region.," Since the x-ray emissivity is proportional to $\rho_{\rm471gas}^2$, the x-ray emission is strongly concentrated into the central region."472 Asstuning the above density profile and the sclfsimilar model as described in Kitavama Suto (1996b). ba typical gamma-ray cluster detectable by the ECRET with AL~WAL. ane 0.05 would have the x-ray flux 2.]«10Heyeeni28 bin 02.1 keV. The inverse-Compton flux is also expected to be conrparable with the thermal enussiou.," Assuming the above density profile and the self-similar model as described in Kitayama Suto (1996b), a typical gamma-ray cluster detectable by the EGRET with $M \sim 10^{15} M_\odot$ and $z \sim4730.05$ would have the x-ray flux $2.4 \times 10^{-11} \rm \ erg \474cm^{-2} \ s^{-1}$ in 0.1–2.4 keV. The inverse-Compton flux is also expected to be comparable with the thermal emission."475 By equating fe: aud faa; du 2.. we eet the cooling photon euergv €-cool=2.0(1|CO keV. below which the electron cooling time is longer than the dvuamical time.," By equating $t_{IC}$ and $t_{\rm shock}$ in \ref{section:flux}, we get the cooling photon energy $\epsilon_{\rm \gamma, cool}476= 2.0 (1+z)^{-5}$ keV, below which the electron cooling time is longer than the dynamical time."477" Then the IC spectrum extends down to around x-ray band withH dN.-/de.XDie.. whileB it becomes harder at waveleusths longer than x-ays with VN.Πε.~e,1."," Then the IC spectrum extends down to around x-ray band with $dN_\gamma/d\epsilon_\gamma \propto \epsilon_\gamma^{-2}$, while it becomes harder at wavelengths longer than x-rays with $dN_\gamma/d\epsilon_\gamma \propto \epsilon_\gamma^{-1.5}$."478 T£ the eamumacrav flux at LOO MeV. is —10©photonsc7s+ that is the EGRET threshold. the IC xav flux (FL) is —1.6«10Hergcni2.41," If the gamma-ray flux at 100 MeV is $\sim 10^{-7} \rm \ photons \ cm^{-2} s^{-1}$ that is the EGRET threshold, the IC x-ray flux $\nu F_\nu$ ) is $\sim 1.6 \times 10^{-11} \rm \ erg \ cm^{-2} \ s^{-1}$."479 Therefore. the thermal and IC fluxes are well above the Bux linüt —bs10Perecmὃν+ of the RBCS.," Therefore, the thermal and IC fluxes are well above the flux limit $\sim 4 \times 10^{-12} \rm480\ erg \ cm^{-2} \ s^{-1}$ of the RBCS."481 However. it takes nearly the ανασα. time for the cluster gas to reach livdrostatic equilibrium after the collapse. aud seanuna-rayvs from the shock generated bv the gravitational collapse are radiated away within that period.," However, it takes nearly the dynamical time for the cluster gas to reach hydrostatic equilibrium after the collapse, and gamma-rays from the shock generated by the gravitational collapse are radiated away within that period."482 Then it is likely that the deusity profile of eanumnia-rav cluitting clusters is more nregular aud extended than ordinary x-ray clusters., Then it is likely that the density profile of gamma-ray emitting clusters is more irregular and extended than ordinary x-ray clusters.483 Iu fact. if the unidentified eur sources in the EGRET catalog are actually extended. they iust have typical angular size of about degree. from the source location accuracy of the EGRET.," In fact, if the unidentified `em' sources in the EGRET catalog are actually extended, they must have typical angular size of about degree, from the source location accuracy of the EGRET."484 As we lave shown. augular size of about 1? is theoretically reasonable if the emission is exteuded to the virial radius.," As we have shown, angular size of about $^\circ$ is theoretically reasonable if the emission is extended to the virial radius."485 When the density profile is not concentrated into the central region but rather constant within the virial radius. the x-ray hinunosity becomes lower than the selfsimilar model by a factor of ~3.7 because of the lower ceutral deusity.," When the density profile is not concentrated into the central region but rather constant within the virial radius, the x-ray luminosity becomes lower than the self-similar model by a factor of $\sim 3.7$ because of the lower central density."486 Furthermore. the surface brieltucss of such loose clusters should be drastically cünuuer than ordinary x-ray clusters.," Furthermore, the surface brightness of such loose clusters should be drastically dimmer than ordinary x-ray clusters."487 In the selfsimular model with ma.2αι the core eas density iS pase~(1/3)GarepassisoDÜPaassis where pan=(Qp/Qulpair is the virial gas density hat is the average gas deusitv within ma.," In the self-similar model with $r_{\rm vir} \gg r_c$, the core gas density is $\rho_{\rm gas, c} \sim (1/3)(r_{\rm vir}/r_{c})^2 488\rho_{\rm gas, vir} \sim 50 \rho_{\rm gas, vir}$, where $\rho_{\rm gas, vir} = (\Omega_B/\Omega_0) \rho_{\rm vir}$ is the virial gas density that is the average gas density within $r_{\rm vir}$."489 Ou the other iud. if the eas density profile of gamma-ray clusters is roughly constant at pai; out to ri. the x-ray surface xiehtuess of such a loose cluster is dinuner than the central surface brightuess of the sclfsimilar model bv a ‘actor of ~(rfras)pssse/Passi)ce200. since the x-rav cluissivity is proportional to Pons," On the other hand, if the gas density profile of gamma-ray clusters is roughly constant at $\rho_{\rm gas, vir}$ out to $r_{\rm vir}$, the x-ray surface brightness of such a loose cluster is dimmer than the central surface brightness of the self-similar model by a factor of $\sim (r_c/r_{\rm vir})(\rho_{\rm gas, c}/490\rho_{\rm gas, vir})^2 \sim 200$, since the x-ray emissivity is proportional to $\rho_{\rm gas}^2$."491 Tt crucially affects he detectability of x-ravs from eamunaray clusters., It crucially affects the detectability of x-rays from gamma-ray clusters.492 The detectability of x-rays should be described by the sigual-o-noise ratio (9/.N) against the x-ray backgrouud fiux, The detectability of x-rays should be described by the signal-to-noise ratio $(S/N)$ against the x-ray background flux.493" The noise level is proportional to (image yee,1/2 and rence SYNxFr where Fo and r are the flux aud the oenage radius. respectively,"," The noise level is proportional to (image $^{1/2}$, and hence $S/N \propto F/r$, where $F$ and $r$ are the flux and the image radius, respectively."494 We have compared the value of Fr of the extended ganuna-ray clusters detectable by the EGRET aud those of the clusters in the RBCS., We have compared the value of $F/r$ of the extended gamma-ray clusters detectable by the EGRET and those of the clusters in the RBCS.495 We fouud that the Fr of gamunarav clusters is by a factor of 3 sunaller than the mium Fr of the RBCS clusters., We found that the $F/r$ of gamma-ray clusters is by a factor of 3 smaller than the minimum $F/r$ of the RBCS clusters.496 The absence of association between the RBCS aud the ECRET sources is therefore not in contracliction to our scenario., The absence of association between the RBCS and the EGRET sources is therefore not in contradiction to our scenario.497 Ou the other haud. deeper observation of candidate gamiuna-rav clusters by Newton. for example. wight detect the x-rav Cluission extended to about 17 with the flux estimated above. that would provide a clear test of our scenario.," On the other hand, deeper observation of candidate gamma-ray clusters by Newton, for example, might detect the x-ray emission extended to about $1^\circ$ with the flux estimated above, that would provide a clear test of our scenario."498 Such x-ray cluission should reflect the structure of shocks in dvuaimically forming clusters. aud nauagiug study is of eyoat interest.," Such x-ray emission should reflect the structure of shocks in dynamically forming clusters, and imaging study is of great interest."499 Tere we again curphasize that the gamma-ray clusters are expected to be more exteuded than clusters that have already stabilized., Here we again emphasize that the gamma-ray clusters are expected to be more extended than clusters that have already stabilized.500" It is known that the surface deusitv profile of galaxies ina cluster cau well be deseribed by the Ising profile. v(r)x[1leer,M lowith the core radius of ~ LOO kpe that is comparable with the core radius of x-ray profile (e.g... Adami et al."," It is known that the surface density profile of galaxies in a cluster can well be described by the King profile, $\sigma(r) \propto [1 +501(r/r_c)^2]^{-1}$ with the core radius of $\sim$ 100 kpc that is comparable with the core radius of x-ray profile (e.g., Adami et al."502 1998)., 1998).503 If we assume a roughly constant surface deusity out to ~rq rather than the Ning profile for gamuna-ray clusters. the average surface deusity," If we assume a roughly constant surface density out to $\sim r_{\rm504vir}$ rather than the King profile for gamma-ray clusters, the average surface density"505"Since the astrometric signature is hidden iu the second derivative of £(2). the computationally most appealing expansion of R aud £ is a Tavlor series around some reference value of :: We define refractivity iutegrals as a short-cut to the notation. covering ((8)) as a special case: A stable nunerical scheme for these iuteerals is proposed in retsec.Ru αν,","Since the astrometric signature is hidden in the second derivative of $L(z)$, the computationally most appealing expansion of $R$ and $L$ is a Taylor series around some reference value of $z$: We define refractivity integrals as a short-cut to the notation, covering \ref{eq.RofnInt}) ) as a special case: A stable numerical scheme for these integrals is proposed in \\ref{sec.Rnum} ."506 Insertion of the series (32)) into the oof (S)) aud iuto the aremments τρτΠο the sines at the vields the expansion cocticicuts with the doublvi-udexed shorthiuds Iu ((33)) C12)) aud refsec.IRuun.. the subscripts of Π are the expoucutial j of the definition (33)): elsewhere they indicate the telescope uumuboer/site.," Insertion of the series \ref{eq.rtayl}) ) into the of \ref{eq.RofnInt}) ) and into the arguments $z_0=z-R$ of the sines at the yields the expansion coefficients with the doubly-indexed shorthands In \ref{eq.RofnInt2}) \ref{eq.Renddef}) ) and \\ref{sec.Rnum}, the subscripts of $R$ are the exponential $j$ of the definition \ref{eq.RofnInt2}) ); elsewhere they indicate the telescope number/site."507 e din ((032)) is of the order of b/(2p) if the reference azimuth + is chosen close to the uiddle between the telescopes. and therefore uot larger than 1.6-10' ad for b«200 un. Because the £j are approximately of the same maguitude refRtavl.ps)). collecting the terms up to j=3 ought establish a relative accuracy of ~5-101 in the anele of refraction.," $x$ in \ref{eq.rtayl}) ) is of the order of $b/(2\rho)$ if the reference azimuth $z$ is chosen close to the middle between the telescopes, and therefore not larger than $1.6\cdot 10^{-5}$ rad for $b<200$ m. Because the $\xi_j$ are approximately of the same magnitude \\ref{Rtayl.ps}) ), collecting the terms up to $j=3$ ought establish a relative accuracy of $\approx 5\cdot 10^{-14}$ in the angle of refraction."508 The expansion proceeds via insertion of ((32)) iuto the sines of ((26)). aud employs an auxiliary. set of iuteerals," The expansion proceeds via insertion of \ref{eq.rtayl}) ) into the sines of \ref{eq.LInt}) ), and employs an auxiliary set of integrals"509For a prolate or oblate ellipsoid we know that where &>1 for a prolate ellipsoid. aud «4<1 for an oblate ellipsoid.,"For a prolate or oblate ellipsoid we know that where $u>1$ for a prolate ellipsoid, and $u<1$ for an oblate ellipsoid."510 So in zx plane iFrom standard trigonometryFrom figure l we can see where n ds the eradieut. aud will uccessarily be the opposite sign to cj.," So in z-x plane >From standard trigonometry>From figure \ref{ellipse} we can see where m is the gradient, and will necessarily be the opposite sign to $x_0$."511 This gives the z value for the z-axis/tangent intercept., This gives the z value for the z-axis/tangent intercept.512 Using (17)) we ect We define the apparcut axial ratio (4) to be less than 1 for both prolate aud oblate ellipsoids., Using \ref{cot}) ) we get We define the apparent axial ratio $q$ ) to be less than 1 for both prolate and oblate ellipsoids.513 Thus for oblate ellipsoids 4=ee and for prolate cllipsoids 4= Z7., Thus for oblate ellipsoids $q=\frac{uA}{a}$ and for prolate ellipsoids $q=\frac{a}{uA}$ .514 Se, So515For many vears there have been theoretical predictions that accretion discs can undergo outbursts which start. either close to their inner edge. or their outer edge (Meyer MaverHofmeister 1984: Mineshige Osaki 1985: Ludwig Mover 1998).,"For many years there have been theoretical predictions that accretion discs can undergo outbursts which start either close to their inner edge, or their outer edge (Meyer Mayer–Hofmeister 1984; Mineshige Osaki 1985; Ludwig Meyer 1998)."516" In this and the following paper we use the eclipsing dwarf novae to resolve spatially their outbursts. and prove that. one. outburst is ""outsidein. the other ""insideout” (Webbetal1999)."," In this and the following paper we use the eclipsing dwarf novae to resolve spatially their outbursts, and prove that one outburst is “outside–in”, the other “inside–out” \cite{Webb99}."517". This is the first time an ""insideoul” outburst has been resolved. anc should be viewed as a major vindication of dise instability. theories."," This is the first time an “inside–out” outburst has been resolved, and should be viewed as a major vindication of disc instability theories."518 1n both cases we find that the discs have extensive vertical structure. which is shown to be the cause of. problems of interpretation in the only other observation of an in” outburst (Vogt. 1983: Rutten 1992).," In both cases we find that the discs have extensive vertical structure, which is shown to be the cause of problems of interpretation in the only other observation of an ``outside--in'' outburst (Vogt 1983; Rutten 1992)."519 The dwarf nova observed in this paper is the SU UAla star IUE Cas. whose orbital period is about LOG min. (," The dwarf nova observed in this paper is the SU UMa star HT Cas, whose orbital period is about 106 min. ("520Patterson 1981: Zhang. Robinson Nather 1986: Wood. llorne Vennes 1992) The observations of LEE Cas. from the 1995. November erüption were carried out with the CCD system. on the 0(.95m James Gregory ‘Telescope (JOUR) at the University of St. Andrews (Bell. Lilditeh Eclwin (1993))). the Thornton Rellector at Ixeele Observatory (Somersctal.1996b) and with a 0.3m Newtonian rellector in Essex.,"Patterson 1981; Zhang, Robinson Nather 1986; Wood, Horne Vennes 1992) The observations of HT Cas, from the 1995 November eruption were carried out with the CCD system on the 0.95m James Gregory Telescope (JGT) at the University of St. Andrews (Bell, Hilditch Edwin \shortcite{Bell93}) ), the 0.6-m Thornton Reflector at Keele Observatory \cite{Somers96b} and with a 0.3m Newtonian reflector in Essex."521 ‘Table 1 lists all the observations., Table \ref{tab:observation_log} lists all the observations.522 Timings are in Darycentric Dynamical Julian Date (BDJD)., Timings are in Barycentric Dynamical Julian Date (BDJD).523 The raw images whicre processed intje standard wav for the instruments in question. as outlines in Somersctal.(1996b) and Bell (1993).," The raw images where processed in the standard way for the instruments in question, as outlined in \scite{Somers96b} and Bell \shortcite{Bell93}."524. The data [rom Essex were dark subtractecl and. [la ficlelecl in the sandard wav [for CCD photometry., The data from Essex were dark subtracted and flat fielded in the standard way for CCD photometry.525 The data were extracted using an optimal extraction technique by avlor (1998)., The data were extracted using an optimal extraction technique by Naylor \shortcite{Naylor98}.526. Eve values for the visual magnitudes for the standard stars were adopted [rom Misselt.(1996)., The values for the visual magnitudes for the standard stars were adopted from \scite{Misselt96}.527. The overall outburst light curve is depicted on Fig. 1.., The overall outburst light curve is depicted on Fig. \ref{fig:alllcvs}.528 The rise to outburst looks very steep with a linear decline., The rise to outburst looks very steep with a linear decline.529 Our coverage of the outburst rise phase starts at a point, Our coverage of the outburst rise phase starts at a point530Following the initial selection of 2>0.9 AMIS8G6 members. stars were rejected if. a) thev are photometric nonmembers inany of our six CMDs. b) they are radial velocity from our spectra or from those spectra used in the Vig study. or ο) they appear to be probable binaries based on a photometric and spectroscopic analyses.,"Following the initial selection of $P>0.9$ MS86 members, stars were rejected if, a) they are photometric nonmembers in of our six CMD's, b) they are radial velocity non-members from our spectra or from those spectra used in the $V_{rad}$ study, or c) they appear to be probable binaries based on a photometric and spectroscopic analyses."531 Figure 1 shows A(Li) as a function of T;jy for the Pleiades. M35 (this study). and the AI35 data from BDSOL are also shown.," Figure 1 shows $A(Li)$ as a function of $T_{eff}$ for the Pleiades, M35 (this study), and the M35 data from BDS01 are also shown."532 All data [vom the literature have been reanalvzed using our methods and photometry. where available. (ο give the most consistent picture possible.," All data from the literature have been reanalyzed using our methods and photometry, where available, to give the most consistent picture possible."533 Several aspects of the 7;;; morphology in M35 are worthy of notice., Several aspects of the $T_{eff}$ morphology in M35 are worthy of notice.534 First. it is clear (hat. as compared to the nominal upper bound plateau between 6300 and 6800Ix with A(Li) = 3.2. many stars between 6000 and 6700 Ix (or bevond) have depleted their surface Li abundances.," First, it is clear that, as compared to the nominal upper bound plateau between 6300 and 6800K with A(Li) = 3.2, many stars between 6000 and 6700 K (or beyond) have depleted their surface Li abundances."535 Although a small amount of Li depletion cannot be ruled. out for the Pleiades. the Li eap hasdefinitely begun to Form in M35.," Although a small amount of Li depletion cannot be ruled out for the Pleiades, the Li gap has begun to form in M35."536 This is clearly illustrated in Figure 2. which shows the continuum-nornmalized spectra in the Li region for four pairs of stars.," This is clearly illustrated in Figure 2, which shows the continuum-normalized spectra in the Li region for four pairs of stars."537 The stars in each pair have similar {ντε and (hus it is not surprising that their Fe I and Ca I line strengths (which depend primarily on 7;;; and intrinsic abundance) are identical: the Li abundances. however. are strikinglv disparate. (," The stars in each pair have similar $T_{eff}$, and thus it is not surprising that their Fe I and Ca I line strengths (which depend primarily on $T_{eff}$ and intrinsic abundance) are identical; the Li abundances, however, are strikingly disparate. ("538Note that where necessary. spectra of slowly rotating stars were artificially. broadened to match their pair.),"Note that where necessary, spectra of slowly rotating stars were artificially broadened to match their pair.)"539 Although the errors listed in Table 1 are internal (see 82). we stress again that svstematic errors due to uncertainties in E(B—V) or T;jy scale do not allect a differential analvsis of stars ad the same νε," Although the errors listed in Table 1 are internal (see 2), we stress again that systematic errors due to uncertainties in $E(B-V)$ or $T_{eff}$ scale do not affect a differential analysis of stars at the same $T_{eff}$."540 Formally. the Li differences are significant at the 5—10e level: the depth of the depletion far exceeds our errors.," Formally, the Li differences are significant at the $5-10\sigma$ level; the depth of the depletion far exceeds our errors."541 The upper aud lower envelope of stars in the gap region of M35 differ by more (han 0.5 dex. which corresponds to to a factor of more than 3 in abundance (and in line strength).," The upper and lower envelope of stars in the gap region of M35 differ by more than 0.5 dex, which corresponds to to a factor of more than 3 in abundance (and in line strength)."542" Thus. M35 illustrates (hat the Li gap begins to form early,"," Thus, M35 illustrates that the Li gap begins to form early."543 second. the Li depletion region is quite wide in Tell.," Second, the Li depletion region is quite wide in Teff."544 The same four pairs illustrate that its width exceeds 700Ix. (from 6000 to 6700IxX2-). which is clearly wider than the canonical llvades Li gap (the region 6550-67001Ix that contains the most extreme Ivacl depletions).," The same four pairs illustrate that its width exceeds 700K (from 6000 to 6700K+), which is clearly wider than the canonical Hyades Li gap (the region 6550-6700K that contains the most extreme Hyad depletions)."545 Bul perhaps this canonical IHvades gap should be viewed as being part of a wider depletion structure: after all. να Li abundances in the region 6300-6550Ix. show Li dispersions of," But perhaps this canonical Hyades gap should be viewed as being part of a wider depletion structure: after all, Hyad Li abundances in the region 6300-6550K show Li dispersions of"546LDNIG621 is a region ofdiffuse emission =25 arcmin to the north of LDNI1622.,LDN1621 is a region of diffuse emission $\approx 25$ arcmin to the north of LDN1622.547 Observations with the CBI at 31 Cllz show a broken ring of emission. that is stronely correlated with PIR emission at 1200100 jim. with Pearson correlation coellicients in the range =0.6 0.5. Optical and Ilo. data show absorption of a strong background. of emission. from warm ionized gas in the Eastern arm of Orion.," Observations with the CBI at 31 GHz show a broken ring of emission, that is strongly correlated with FIR emission at $12-100~\mu$ m, with Pearson correlation coefficients in the range $\approx 0.6-0.8$ Optical and $\alpha$ data show absorption of a strong background of emission from warm ionized gas in the Eastern arm of Orion."548 This suggests that LDNI621 and LDN1622 are in the foreground of Orion (al à distance un500 parsec). possibly as close as 120 pc (Wilsonetal.2005).," This suggests that LDN1621 and LDN1622 are in the foreground of Orion (at a distance of $\sim 500$ parsec), possibly as close as 120 pc \citep{Wilson05}."549.. No Lla emission. associated. clürectlv with LDNIG2I. aseen.," No $\alpha$ emission, associated directly with LDN1621, is seen."550 “This suggests that LDNI621 itself is not emitting significant [rec-free. emission. although the ellects of dust. extinction do not allow a strong constraint to be placed.," This suggests that LDN1621 itself is not emitting significant free-free emission, although the effects of dust extinction do not allow a strong constraint to be placed."551 Low frequency. radio data also do not show evidence of dilfuse emission associated with LDNIG621., Low frequency radio data also do not show evidence of diffuse emission associated with LDN1621.552 The 31 CGllz emission is at z20—30 mJy while an analvsis of the GBG6 map at 4.85 CGllz provides a strong (30) upper limit of 7.2 mJv Lau 31 Guz for frec-[ree emission., The 31 GHz emission is at $\approx 20-30$ mJy $^{-1}$ while an analysis of the GB6 map at 4.85 GHz provides a strong $3\sigma$ ) upper limit of 7.2 mJy $^{-1}$ at 31 GHz for free-free emission.553 The FlR-corrclated emission at 31 Gllz therefore appears to be mostly due to radiation associated with dust., The FIR-correlated emission at 31 GHz therefore appears to be mostly due to radiation associated with dust.554 IRAS data alone do not allow a reliable extrapolation of the RayleighJeans thermal dust tail to 31 Gilz., IRAS data alone do not allow a reliable extrapolation of the Rayleigh-Jeans thermal dust tail to 31 GHz.555WAZAP data at 95 mGllz combined with HUXS data allowed the Dux density to estimatedin an aperture of diameter 30 arcmin at an angular resolution of 13 arcmin., data at 93.5 GHz combined with IRAS data allowed the flux density to be estimated in an aperture of diameter 30 arcmin at an angular resolution of 13 arcmin.556 A single moclified blackbody indicates that the thermal dust is 10 per cent ofthe total 31 CLIz us. corresponding to an excess of 1.524 Jv (2.30).," A single modified blackbody indicates that the thermal dust is $\sim 10$ per cent of the total 31 GHz flux, corresponding to an excess of $1.52\pm0.66$ Jy $2.3\sigma$ )."557 The dust-correlated emission has a coupling coelIicient. relative to LOO pim. of 18.12354.4 pls |. consistent. with that observed from LDN1622.," The dust-correlated emission has a coupling coefficient, relative to $100~\mu$ m, of $18.1\pm4.4~\mu$ K $^{-1}$, consistent with that observed from LDN1622."558 Orion East(consisting of both LDNI621 and LDN1622) appear to be part of the same svstem of dust clouds. emitting significant anomalous emission at [requencies ~30 Cillz.," Orion East (consisting of both LDN1621 and LDN1622) appear to be part of the same system of dust clouds, emitting significant anomalous emission at frequencies $\sim 30$ GHz."559 Spinning dust is an obvious candidate for the physical mechanism responsible for the bulk of the emission., Spinning dust is an obvious candidate for the physical mechanism responsible for the bulk of the emission.560 Ligh sensitivity data. covering a wide range of frequencies. (~5300 Cllz). is required to study such clouds in. more detail.," High sensitivity data, covering a wide range of frequencies $\sim5-300$ GHz), is required to study such clouds in more detail."561 Data from thePlanck satellite will be particularly useful in constraining the Iavleigh-Jeans dust tail. which niv be responsible for a significant fraction of the 31 12 if the emissivity index llattens at longer wavelengths.," Data from the satellite will be particularly useful in constraining the Rayleigh-Jeans dust tail, which may be responsible for a significant fraction of the 31 GHz if the emissivity index flattens at longer wavelengths."562 This work was supported by the Strategic Alliance for the Implementation. of New Technologies. (SAINT. - see wwwsastro.caltech.edu/chajnantorsaint/index.html) and we are most grateful to the SAINT partners for their strong support., This work was supported by the Strategic Alliance for the Implementation of New Technologies (SAINT - see www.astro.caltech.edu/chajnantor/saint/index.html) and we are most grateful to the SAINT partners for their strong support.563 We gratefully acknowledge support from the Ixavli Operating Institute and thank D. Rawn and S. Rawn Jr. Phe CBL was supported by NSE grants 9802089. 0008734 anc 0206416. and a Roval Society Small Research rant.," We gratefully acknowledge support from the Kavli Operating Institute and thank B. Rawn and S. Rawn Jr. The CBI was supported by NSF grants 9802989, 0098734 and 0206416, and a Royal Society Small Research Grant."564 We are particularly indebted to the engineers who maintainec and operated the CBL: Cristobbal Achermann. José Cortéss. Cristóbbal Jara. Nolberto Ovarace. Martin. Shepherd anc Carlos Verdugo.," We are particularly indebted to the engineers who maintained and operated the CBI: Cristóbbal Achermann, José Cortéss, Cristóbbal Jara, Nolberto Oyarace, Martin Shepherd and Carlos Verdugo."565 CD acknowledges an SPEC Xdvanceec Fellowship and ERC erant. under the PPT., CD acknowledges an STFC Advanced Fellowship and ERC grant under the FP7.566 We acknowledge the use of the Legacy Archive for Microwave. Dackgrounc Data Analvsis (LAAIBDA)., We acknowledge the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA).567 Support. for LAAIBDA is provided by the NASA Ollice of Space Science., Support for LAMBDA is provided by the NASA Office of Space Science.568We used data from the Southern L-Alpha Sky Survey Atlas (SLLASSA). which is supported by the National Science Foundation.,"We used data from the Southern H-Alpha Sky Survey Atlas (SHASSA), which is supported by the National Science Foundation."569MLC was supported by theCajal Programme of the Spanish science ministry.,MLC was supported by the Programme of the Spanish science ministry.570 We have used (he following online databases: Sloan Digital Sky Survey (http://www.sdss.org/). ihe Atomic Line List (http://www.pa.ukv.edu/ peter/atomic/). and NIST Atomic Spectra Database (http://physies.nist.gov/PhvsBRefData/ASD).," We have used the following online databases: Sloan Digital Sky Survey (http://www.sdss.org/), the Atomic Line List (http://www.pa.uky.edu/ peter/atomic/), and NIST Atomic Spectra Database (http://physics.nist.gov/PhysRefData/ASD)."571Smoothec Particle Hvdrodynamiecs (SPLHI) was firs introduced. as a tool for studying stellar structure (7. 2: 7 7)). but has since found. wide application in all areas of theoretical astrophysics (?).. in engineering (2).. and bevonc (c.g. ? ?)).,"Smoothed Particle Hydrodynamics (SPH) was first introduced as a tool for studying stellar structure \bcite{1977MNRAS.181..375G}; ; \bcite{1977AJ.....82.1013L}) ), but has since found wide application in all areas of theoretical astrophysics \citep{1992ARA&A..30..543M}, , in engineering \citep{1993JCoPh.109...67L}, , and beyond (e.g. \bcite{2008JCoPh.227.9195H}) )."572 Although there are many varieties of SPLIT. the centra idea is to represent a fIuid. by discrete particles that move with the Dow (?2: 7 2)).," Although there are many varieties of SPH, the central idea is to represent a fluid by discrete particles that move with the flow \bcite{1992ARA&A..30..543M}; \bcite{2005astro.ph..7472P}) )."573 Dvpically these particles represen the [uid exactly. though in some variants the Bui is advected: on top of the particles (27: ? 73).," Typically these particles represent the fluid exactly, though in some variants the fluid is advected on top of the particles \bcite{1999Dilts}; \bcite{2003ApJ...595..564M}) )."574 Phe kev advantages over Eulerian are its Lagrangian nature that makes it Galilean invariant. and its particle nature that makes it easy to couple to the fast multipole method for gravity that scales as O(N) (22: 7? 2)).," The key advantages over Eulerian are its Lagrangian nature that makes it Galilean invariant, and its particle nature that makes it easy to couple to the fast multipole method for gravity that scales as $O(N)$ \bcite{2000ApJ...536L..39D}; \bcite{1987JCoPh..73..325G}) )."575 However. SPL has problems correctly integrating ΕΙ instabilities ancl mixing at boundariesTu YTTu Tu PP 7 PV).," However, SPH has problems correctly integrating fluid instabilities and mixing at boundaries; \bcite{1999Dilts}; \bcite{2001MNRAS.323..743R}; \bcite{2003MNRAS.345..561M}; ; \bcite{2006astro.ph.10051A}) )."576 Several different reasons have been sugeested for this in the literature so far., Several different reasons have been suggested for this in the literature so far.577 7 and ? argue that the problem owes to errors in the SPL gradients that do not show good. convergence for irregular. particle distributions., \citet{1996PhDMorris} and \citet{1999Dilts} argue that the problem owes to errors in the SPH gradients that do not show good convergence for irregular particle distributions.578 ? argue that the problem. owes to the fact that entropies are discontinuous at boundaries. while the densities are smooth.," \citet{2007arXiv0709.2772P} argue that the problem owes to the fact that entropies are discontinuous at boundaries, while the densities are smooth."579 This gives spurious pressure blips at boundaries that drive Duids of dillerent entropy. apart., This gives spurious pressure blips at boundaries that drive fluids of different entropy apart.580 ον find that adding thermal conductivity at. boundaries to smooth the entropies gives improved mixing in SP, They find that adding thermal conductivity at boundaries to smooth the entropies gives improved mixing in SPH.581ILL. ? make a similar argument. phrasing the problem in terms of an inability for SPILL particles to mix ancl generate entropy on the kernel scale.," \citet{2008MNRAS.387..427W} make a similar argument, phrasing the problem in terms of an inability for SPH particles to mix and generate entropy on the kernel scale."582 They. find that adding a heat diffusion term to model subericl turbulence gives improved. mixing in SPI., They find that adding a heat diffusion term to model subgrid turbulence gives improved mixing in SPH.583 Finally. 2. suggest that the problem lies in the SPL density estimate.," Finally, \citet{2001MNRAS.323..743R} suggest that the problem lies in the SPH density estimate."584 They. introduce a new temperature weighted density estimate that is designed to give smoother pressures at. [low boundaries. thus combating the spurious boundary. pressure blip.," They introduce a new temperature weighted density estimate that is designed to give smoother pressures at flow boundaries, thus combating the spurious boundary pressure blip."585 In this paper. we perform an error ancl stability analysis Of SPILL in its most general form to understand why mixing fails.," In this paper, we perform an error and stability analysis of SPH in its most general form to understand why mixing fails."586 In doingthis. weshowthat all of the above authors correctly identified one of twodistinet problems with mixing," In doingthis, weshowthat all of the above authors correctly identified one of twodistinct problems with mixing"587"the initial mass fractions for!?C,,'*N,, and the alpha elements, and a look-up table for the true data block.","the initial mass fractions for, and the alpha elements, and a look-up table for the true data block."588" The final file consists of 63 rectangular data arrays, where logκι is tabulated as a function of logT and logR."," The final file consists of 63 rectangular data arrays, where $\log \kappa_\mathrm{R}$ is tabulated as a function of $\log T$ and $\log R$."589 The tables are ordered such that the mass fraction Χ(12Ο) varies the most rapidly followed by the hydrogen mass fraction and X(!N)., The tables are ordered such that the mass fraction $X(\mbox{\element[][12]{C}})$ varies the most rapidly followed by the hydrogen mass fraction and $X(\mbox{\element[][14]{N}})$.590" For future compatibility, a data field for the alpha element enhancement factor was introduced into the look-up table."," For future compatibility, a data field for the alpha element enhancement factor was introduced into the look-up table."591 We compare our tables based on a scaled solar metal mixture with data from F05 based on the same abundances as in this work., We compare our tables based on a scaled solar metal mixture with data from F05 based on the same abundances as in this work.592" A direct comparison with AF94 is not possible because there are, of course, no tables based on the Lodders(2003) abundances."," A direct comparison with AF94 is not possible because there are, of course, no tables based on the \citet{2003ApJ...591.1220L} abundances."593 We refer to F05 for a comparison of AF94 and FOS., We refer to F05 for a comparison of AF94 and F05.594" In the figures, we always depict data from our database as Kcoma, while the respective comparison values are labelled κκ."," In the figures, we always depict data from our database as $\kappa_\mathrm{COMA}$ , while the respective comparison values are labelled $\kappa_\mathrm{R}$."595" Despite the numerous differences between the COMA and F05 approach, we find reasonable agreement between both sets of data."," Despite the numerous differences between the COMA and F05 approach, we find reasonable agreement between both sets of data."596 For the case shown in Fig., For the case shown in Fig.597" 4 (Z= 0.02), the difference between the COMA and ΕΟ5 values does not exceed 15 per cent for temperatures as low as logT=3.5."," \ref{fig:coma-f05-full-sc} $Z=0.02$ ), the difference between the COMA and F05 values does not exceed 15 per cent for temperatures as low as $\log T=3.5$."598 The discrepancies at lower temperatures are higher (up to 35 per cent) and can in fact be ascribed to several things., The discrepancies at lower temperatures are higher (up to 35 per cent) and can in fact be ascribed to several things.599" First and foremost, the use of different sets of molecular data in the calculations (cf."," First and foremost, the use of different sets of molecular data in the calculations (cf."600 our Table 2 and their Tables 3 and 4) produces a deviation in the resulting mean opacity coefficients., our Table \ref{table:molecules} and their Tables 3 and 4) produces a deviation in the resulting mean opacity coefficients.601" Second, we adopt a microturbulent velocity of 2.5kms!, while FOS use 20kms!."," Second, we adopt a microturbulent velocity of $2.5\,\mathrm{km\,s^{-1}}$, while F05 use $2.0\,\mathrm{km\,s^{-1}}$."602" The choices for this parameter are (within a certain range that is found for atmospheres of low mass giants) somewhat arbitrary and cause perceptible changes in Kg, especially at lower temperatures."," The choices for this parameter are (within a certain range that is found for atmospheres of low mass giants) somewhat arbitrary and cause perceptible changes in $\kappa_\mathrm{R}$, especially at lower temperatures."603" Third, F05 use a denser wavelength grid for the evaluation of xg."," Third, F05 use a denser wavelength grid for the evaluation of $\kappa_\mathrm{R}$ ."604 We discuss these issues in more detail in Sect. ??.., We discuss these issues in more detail in Sect. \ref{sec:uncertainties}.605 From a comparison of Fig., From a comparison of Fig.606 4 with Figs., \ref{fig:coma-f05-full-sc} with Figs.607" 10 (showing a comparable order of magnitude of the deviations) and 11, it is, however, clear that the numerous differences in the physical input data are responsible for the major part of the discrepancies."," \ref{fig:coma-h2o-c2-relative-logT} (showing a comparable order of magnitude of the deviations) and \ref{fig:coma-xi-f05res-relative-logT}, it is, however, clear that the numerous differences in the physical input data are responsible for the major part of the discrepancies."608 The resolution and microturbulent velocity influence kg not quite as much., The resolution and microturbulent velocity influence $\kappa_\mathrm{R}$ not quite as much.609" The large deviations in the data at the lowest temperatures are due to grain opacity that we do not take into account in our calculations, but dust is usually not formed under equilibrium conditions (as assumed by F05, see Introduction)."," The large deviations in the data at the lowest temperatures are due to grain opacity that we do not take into account in our calculations, but dust is usually not formed under equilibrium conditions (as assumed by F05, see Introduction)."610" Moreover, FOS adopted a finer grid in logT below 3.5."," Moreover, F05 adopted a finer grid in $\log T$ below 3.5."611" For the oxygen-rich case, a cubic spline interpolation (see Fig. 4,,"," For the oxygen-rich case, a cubic spline interpolation (see Fig. \ref{fig:coma-f05-full-sc},"612 dotted lines) on the coarser grid we adopted (and also used by AF94) provides reasonably accurate values., dotted lines) on the coarser grid we adopted (and also used by AF94) provides reasonably accurate values.613 The comparison with high temperature data such as that from OPAL or OP is limited to the temperature regions where the tables overlap., The comparison with high temperature data such as that from OPAL or OP is limited to the temperature regions where the tables overlap.614" Moreover, it is this region where a transition between low and high temperature opacities has to be made for applications covering a wide temperature range."," Moreover, it is this region where a transition between low and high temperature opacities has to be made for applications covering a wide temperature range."615" OP data stretch down to logT3.5, whereas the OPAL tables end at logT3.75."," OP data stretch down to $\log T=3.5$, whereas the OPAL tables end at $\log T=3.75$."616 The comparison for a standard scaled solar composition in Fig., The comparison for a standard scaled solar composition in Fig.617 5 shows a growing deviation for lower temperatures because both OPAL and OP do not include molecular absorbers (except H5)., \ref{fig:coma-op-opal-full} shows a growing deviation for lower temperatures because both OPAL and OP do not include molecular absorbers (except $_2$ ).618" This plot indicates that in the region between logT=3.8 and the high temperature end of the COMA data, a smooth transition to high temperature data is possible."," This plot indicates that in the region between $\log T=3.8$ and the high temperature end of the COMA data, a smooth transition to high temperature data is possible."619" Again, from the dimension of the differences, we conclude that these are due to different physical input data ratherthan other parameters (see Sect. ??))."," Again, from the dimension of the differences, we conclude that these are due to different physical input data ratherthan other parameters (see Sect. \ref{sec:uncertainties}) )."620 To assess which temperature region renders itself tosuch a crossover we plot the logarithmic difference between our, To assess which temperature region renders itself tosuch a crossover we plot the logarithmic difference between our621Because of very effective tidal truncation on Be disks in the relatively narrow and low eccentric svstems. the flow of the matter towards the DII will be effectively blocked during almost the whole orbital evele.,"Because of very effective tidal truncation on Be disks in the relatively narrow and low eccentric systems, the flow of the matter towards the BH will be effectively blocked during almost the whole orbital cycle."622 In such episode. the accretion could be mainly from the polar wind of the donor star.," In such episode, the accretion could be mainly from the polar wind of the donor star."623 As suggested by Waters&vanNerkwijk(1989).. the polar wind of the Be star probably resembles that in OD stars.," As suggested by \citet{wat89}, the polar wind of the Be star probably resembles that in OB stars."624 Such low-clensity high-velocity wind can hardly form accretion disks around the BIIs. and the accretion would follow the classical Bondi-IHovle-Littleton. (BILL) approximation.," Such low-density high-velocity wind can hardly form accretion disks around the BHs, and the accretion would follow the classical Bondi-Hoyle-Littleton (BHL) approximation."625" For a typical main-sequence 15AL. main-sequence star in the svstem of 2,=10d. if the mass loss rate M.—10""M.vr! and the terminal wind velocity ος=2000kms+ exist. the accreting compact star would have luminosity of about L;~LOeres+."," For a typical main-sequence $15 \, M_{\sun}$ main-sequence star in the system of $\Porb=10 \, {\rm d}$, if the mass loss rate $\Mwdot=10^{-9} \, M_{\sun}\ {\rm yr}^{-1}$ and the terminal wind velocity $v_{\infty}=2000 \, {\rm km \, s}^{-1}$ exist, the accreting compact star would have luminosity of about $\Lx\sim62610^{33} \, \rm{erg \, s^{-1}}$."627 For the De/DII binaries. the luminosity would be still less. since the polar wind just takes a small fraction of the mass loss from the donor star.," For the Be/BH binaries, the luminosity would be still less, since the polar wind just takes a small fraction of the mass loss from the donor star."628 The wind plasma accumulates in the outer rings of the decretion disc 41 one-armed oscillaàton instabilitv - probably responsible for the V/I variability seen in the Be stars - occurs. Which causes the almost total disruption of the Be disk.," The wind plasma accumulates in the outer rings of the decretion disc till one-armed oscillaiton instability - probably responsible for the V/R variability seen in the Be stars - occurs, which causes the almost total disruption of the Be disk."629 The resulted. large onto the black hole would lead (o very. luminous outbursts., The resulted large mass-infall onto the black hole would lead to very luminous outbursts.630 However. no aspect of the truncation model implies (hat the large mass transfer must occur (Okazakiοἱal.2002).," However, no aspect of the truncation model implies that the large mass transfer must occur \citep{oka02}."631.. II the collapsed disk due to the ανπάσα} instability falls back on to the more massive Be star. such narrow and low-eccentricitv transients would be almost out. of detection even lor the most sensitive instruments. so it should be careful to deal with the Be stars unrelated to the X-raw SOTIECES.," If the collapsed disk due to the dynamical instability falls back on to the more massive Be star, such narrow and low-eccentricity transients would be almost out of detection even for the most sensitive instruments, so it should be careful to deal with the Be stars unrelated to the X-ray sources."632 For systems in which effective (uneation occurs. e.g. NTE J1543—568 and 25 1553—542. the burst activities are very rare (see Okazakietal.(2002) ancl reference therein).," For systems in which effective truncation occurs, e.g. XTE $1543-568$ and 2S $1553-542$, the burst activities are very rare (see \citet{oka02} and reference therein)."633 During the long quiescent stage. De stars could be identified [rom (he Balmer (and sometimes other) line emission and the associated strong infrared excess.," During the long quiescent stage, Be stars could be identified from the Balmer (and sometimes other) line emission and the associated strong infrared excess."634 In the IIMXD catalogue of Liu.ParaclijsHeuvel (2000).. there are probably 24 De/X-ray systems where the nature of the compact star and (he orbital period are undetermined.," In the HMXB catalogue of \citet{liu00}, there are probably 24 Be/X-ray systems where the nature of the compact star and the orbital period are undetermined."635 They are signed as ΠοΝταν binaries because of their highly variable X-ray. characteristics analogous to those of the well studied Be/NS transients. or the optical identification of the donor stars.," They are signed as Be/X-ray binaries because of their highly variable X-ray characteristics analogous to those of the well studied Be/NS transients, or the optical identification of the donor stars."636 Most of them have been observed once during the burst states. and (hen disappeared [roi X-ray detection because ol verv low Iuminositv.," Most of them have been observed once during the burst states, and then disappeared from X-ray detection because of very low luminosity."637 Among these undetermined De/X-rav svstems. NTE J1739—302 and AX J0052.9—7158 (SAIC 32) can be excluded now since the donor star of the former has been identified as an O supereiant (Smithetal.2003) ancl 167.8 seconds pulsations have been found in the latter (Yokogawaetal.2001).," Among these undetermined Be/X-ray systems, XTE $1739-302$ and AX $0052.9-7158$ (SMC 32) can be excluded now since the donor star of the former has been identified as an O supergiant \citep{smi03}638 and 167.8 seconds pulsations have been found in the latter \citep{yok01}."639. To look for Be/DII binaries in the left svstenis. optical observations in the quiescent stage become important. from which we can confinn the nature of the donor star and get information of the orbital period aud (the velocity. curve of the donor.," To look for Be/BH binaries in the left systems, optical observations in the quiescent stage become important, from which we can confirm the nature of the donor star and get information of the orbital period and the velocity curve of the donor."640 The latter (wo are useful in estimating the dynamical mass of the compact star - (he common wav to determine whether the compact star is a black hole. though the," The latter two are useful in estimating the dynamical mass of the compact star - the common way to determine whether the compact star is a black hole, though the"641the beeinning of a evele and then progressively move to lower latitudes.,the beginning of a cycle and then progressively move to lower latitudes.642 This plot looks verv much like the plots presented by Durnev (1997) especiallv his Figure 7., This plot looks very much like the plots presented by Durney (1997) — especially his Figure 7.643 This is cerlainly very reassuring. since (he numerical techniques emploved by us aud by Durney (1997) are completely different.," This is certainly very reassuring, since the numerical techniques employed by us and by Durney (1997) are completely different."644 Apart from the production of the double rings. our code allows lor the toroidal flux to be brought to the surface bv meridional circulation and then to be acted upon by a-coefficient (an effect not present in Durneys calculations).," Apart from the production of the double rings, our code allows for the toroidal flux to be brought to the surface by meridional circulation and then to be acted upon by $\alpha$ -coefficient (an effect not present in Durney's calculations)."645" However. when A""is made as large as 1000. this effect is insignificant."," However, when $K'$ is made as large as 1000, this effect is insignificant."646 In fact. we made some runs with a=0 and found that the results for zero or non-zero a are virtually indistinguishable when A’=1000.," In fact, we made some runs with $\alpha = 0$ and found that the results for zero or non-zero $\alpha$ are virtually indistinguishable when $K' = 1000$."647 For example. the plots of eruption latitude against time ancl the butterfly diagrams look identical in both the cases.," For example, the plots of eruption latitude against time and the butterfly diagrams look identical in both the cases."648" We have already mentioned that a positive A"" is like a positive a-elleet concentrated near the surface.", We have already mentioned that a positive $K'$ is like a positive $\alpha$ -effect concentrated near the surface.649 Choudhunr. Sehüsssler. Dikpati (1995) showed that a positive a concentrated near the surface leads to a poleward propagation of the dynamo wave when the meridional flow is switched off.," Choudhuri, Schüsssler, Dikpati (1995) showed that a positive $\alpha$ concentrated near the surface leads to a poleward propagation of the dynamo wave when the meridional flow is switched off."650" We find exactly the same result in the double ring approach with positive A"" if we switch off the meridional flow.", We find exactly the same result in the double ring approach with positive $K'$ if we switch off the meridional flow.651 Figure 5 shows a time-latitude plot of the toroidal field al the bottom of the convection zone with meridional flow for the case A!=1000. fy=0.5. whereas Figure 6 is a similar plot without meridional flow. keeping all the other parameters (he same.," Figure 5 shows a time-latitude plot of the toroidal field at the bottom of the convection zone with meridional flow for the case $K'=1000$, $f_d = 0.5$, whereas Figure 6 is a similar plot without meridional flow keeping all the other parameters the same."652 We see clear indication of poleward migration in Figure 6., We see clear indication of poleward migration in Figure 6.653 For contrast. we now present results obtained by the method described in 2.2.," For contrast, we now present results obtained by the method described in 2.2."654 As we have seen. the control parameter in (his problem is /(«1). which measures the strength of magnetic buovaneyv.," As we have seen, the control parameter in this problem is $f (<1)$, which measures the strength of magnetic buoyancy."655 Figure 7 shows how the dvnamo period changes on increasing f., Figure 7 shows how the dynamo period changes on increasing $f$.656 As in Figure 2. we begin with a period of 66 vrs in the limit /=0 corresponding to the CSD model.," As in Figure 2, we begin with a period of 66 yrs in the limit $f= 0$ corresponding to the CSD model."657 On making the effect of buovancy stronger (by increasing |). the flix transport," On making the effect of buoyancy stronger (by increasing $f$ ), the flux transport"658those of Seyferts and. QSOs.,those of Seyferts and QSOs.659" Lowe consider the ""wash out” fact for LLAGNs. the correlation should be tighter."," If we consider the “wash out” fact for LLAGNs, the correlation should be tighter."660 One may notice that NGC3227 deviates from the main trend in Figure L.., One may notice that NGC3227 deviates from the main trend in Figure \ref{fig-1}.661 This object is unusual by virtue of its very [at soft. X-ray spectrum which may. be caused by a clusty warm absorber (Georecetal.1998:Ixomassa&Fink 1991).," This object is unusual by virtue of its very flat soft X-ray spectrum which may be caused by a dusty warm absorber \cite{george,kf}."662. Probably its variability is enhanced by some changes in the absorber., Probably its variability is enhanced by some changes in the absorber.663 In addition. Schinnerer. Eckart Tacconi (2000) reported that the enclosed mass in the inner 25 pe of NGC3221 is about 210AZ. based on a detection of molecular gas at a clistance from nucleus of only ~15 pc.," In addition, Schinnerer, Eckart Tacconi \shortcite{set} reported that the enclosed mass in the inner 25 pc of NGC3227 is about $2\times10^7M_{664\odot}$ based on a detection of molecular gas at a distance from nucleus of only $\sim15$ pc."665 Although this mass approximately agrees with the Virial mass of 3.94.0«10AL; measured by using reverberation mapping cata with larger uncertainty. it may suggest that the mass of the central black hole is lower than the measured Virial mass.," Although this mass approximately agrees with the Virial mass of $3.9-4.9\times10^7M_{\sun}$ measured by using reverberation mapping data with larger uncertainty, it may suggest that the mass of the central black hole is lower than the measured Virial mass."666 Wothe central black hole mass of NGC3227 is really lower than the estimated. Virial mass (sav. by a factor of 2). NGC€C3227 would join the main trend.," If the central black hole mass of NGC3227 is really lower than the estimated Virial mass (say, by a factor of 2), NGC3227 would join the main trend."667 Furthermore. we ga10uld caution that some systematic error may exist because je adopted masses are measured. by cilferent techniques for LLAGNs and Sevfert 1 ealaxies/OQSOs.," Furthermore, we should caution that some systematic error may exist because the adopted masses are measured by different techniques for LLAGNs and Seyfert 1 galaxies/QSOs."668 This svstematic error should not be very large since the measured masses by 1ese two methods follow the same relation with the galaxies rulge potential (Gebhardtctal.2000b:Nelson2000).. and 1 trend still remains even i we exclude those points for LLAGNs in Figure 1..," This systematic error should not be very large since the measured masses by these two methods follow the same relation with the galaxies bulge potential \cite{g2000b,nelson}, and the trend still remains even if we exclude those points for LLAGNs in Figure \ref{fig-1}."669" 1n Figure L.. the trend seems to be the case that there is. a linear. relationship.. στιµ..2xMdb,"," In Figure \ref{fig-1}, the trend seems to be the case that there is a linear relationship, $\sigma^2_{\rm rms}\propto M^{-1}_{\rm bh}$."670" As we can sec. it can be represented by a line logez,Lus435log(Myun/M.) in both Figure 1 A and. B. All the objects are localized. in the region between the lines logaz.=3.75log(AlinM.) and loge,2—5.55log(Ady/M.) except NGC3227."," As we can see, it can be represented by a line $\log\sigma^2_{\rm rms}=4.75-\log(671M_{\rm bh}/M_{\odot})$ in both Figure \ref{fig-1} A and B. All the objects are localized in the region between the lines $\log\sigma^2_{\rm 672rms}=3.75-\log(M_{\rm bh}/M_{\odot})$ and $\log\sigma^2_{\rm rms}=5.75-\log(M_{\rm bh}/M_{\odot})$ except NGC3227."673" More quantitatively, a Spearman rank test gives the correlation coelficients of -0.70 and -0.65 for the points in Figure 1 A and D. respectively: and rejects the possibility that στης and Ady, are uncorrelated at 2299.94 confidence."," More quantitatively, a Spearman rank test gives the correlation coefficients of -0.70 and -0.65 for the points in Figure \ref{fig-1} A and B, respectively; and rejects the possibility that $\sigma^2_{\rm rms}$ and $M_{\rm bh}$ are uncorrelated at $>$ confidence."674 Vhe robust nature of a rank test means that the significance of this correlation does not depend on the outlying point λος2291., The robust nature of a rank test means that the significance of this correlation does not depend on the outlying point NGC3227.675 Η we include NCGC3227 in the test. the corresponding correlation coefficient are -0.68 and. -0.64 with almost the same confidence for the points in Figure 1. X and 0. respectively.," If we include NGC3227 in the test, the corresponding correlation coefficient are -0.68 and -0.64 with almost the same confidence for the points in Figure \ref{fig-1} A and B, respectively."676" In the present paper. we found that the ""excess variance"" is significantly anti-correlated. with the central black hole mass for a combined sample of Sevfert 1 galaxies. QSOs anc LLAGNs."," In the present paper, we found that the “excess variance” is significantly anti-correlated with the central black hole mass for a combined sample of Seyfert 1 galaxies, QSOs and LLAGNs."677 The most plausible explanation is that the “excess variance is caused by some global coherent changes in the X-ray emitting region. and this region scales with the size of black hole.," The most plausible explanation is that the “excess variance” is caused by some global coherent changes in the X-ray emitting region, and this region scales with the size of black hole."678" Phe light curves are known to be characterize by à steep power-Iaw PDS (PCP)xf"".wheree —L5.2) in some ACGNs. such as NG€4051. NGC3516. NGC€CB5AS ane ALCC-6-30-15 (Lawrence&Papaclakis1993:NowakChi-ang2000)."," The light curves are known to be characterized by a steep power-law PDS $P(f)\propto f^{-\alpha}$, where $\alpha679\sim 1.5-2$ ) in some AGNs, such as NGC4051, NGC3516, NGC5548 and MCG-6-30-15 \cite{lp,nch}."680". Assuming self similar scaling and hence a direc connection between time scales and the size of sources. the observed. ""excess variance” can be related to the size of the central black hole as στ.=IDPODdfxfi©Al. ""s where fi; and az1."," Assuming self similar scaling and hence a direct connection between time scales and the size of sources, the observed “excess variance” can be related to the size of the central black hole as $\sigma^2_{\rm rms} = \int_{f_1}^{f_2}P(f)df \propto f^{1-\alpha}_{1}681\propto R^{1-\alpha} \propto M^{1-\alpha}_{\rm bh}$ , where $f_1\ll f_2$ and $\alpha\neq 1$."682" One can readily get the observed. correlation ez,xAd, lcillustrated in. Figurem 1 by assuming à~2."," One can readily get the observed correlation $683\sigma^2_{\rm rms} \propto M^{-1}_{\rm bh}$ illustrated in Figure \ref{fig-1}684 by assuming $\alpha \sim 2$."685" This fundamental relationship can self-consistently explain the previous finding of the relationship between oz, and Luminosity. which has been proposed.by many authors (Nandraetal.LOOT:Turner.1999:Leighlv1999:Xlmainietal."," This fundamental relationship can self-consistently explain the previous finding of the relationship between $\sigma^2_{\rm rms}$ and luminosity, which has been proposedby many authors \cite{nandra,turner,leighly,alm}."686" 2000).. ltecentIs. some investigations have been performed. on the relationship between of, and Luminosity for a sample of LLAGNs (Ptakctal.1998). anc a deep flux. limited sample of QSOs selected (rom deep ROSAT survey CAlmainietal. 2000). which greatly extend the luminosity range and redshift range."," Recently, some investigations have been performed on the relationship between $\sigma^2_{\rm rms}$ and luminosity for a sample of LLAGNs \cite{ptak}687 and a deep flux limited sample of QSOs selected from deep ROSAT survey \cite{alm}, which greatly extend the luminosity range and redshift range."688 Ptak et al., Ptak et al.689 (1998). found. that LLAGNs tend to show little or no significant short termi variability. and there is a break from the trend of increased variability in Sevfert 1 galaxies with decreased Luminosity.," \shortcite{ptak} found that LLAGNs tend to show little or no significant short term variability, and there is a break from the trend of increased variability in Seyfert 1 galaxies with decreased luminosity."690 They proposed that this is cue to the lower aceretion rates in LLACGNs., They proposed that this is due to the lower accretion rates in LLAGNs.691 Vhey arguedὃν that this results in a largero characteristic size of the X-ray emission region in LLAGNs than in Sevfert 1 ealaxies because the lower accretion rate is probably causing the accretion Llow to be advection-dominated., They argued that this results in a larger characteristic size of the X-ray emission region in LLAGNs than in Seyfert 1 galaxies because the lower accretion rate is probably causing the accretion flow to be advection-dominated.692 However. most of the X-ray emission should originate in an inner volume probably with a radius less than 10/22: (Ptakοἱal. 1998).. which is similar to the typical X-ray emission region size ( LORS) of a normal Sevfert. 1 galaxies.," However, most of the X-ray emission should originate in an inner volume probably with a radius less than $R_{\rm Sch}$ \cite{ptak}, which is similar to the typical X-ray emission region size $\sim 10R_{\rm Sch}$ ) of a normal Seyfert 1 galaxies."693 Lf the X-ray variability is caused by some global coherent oscillation for both LLAGNs and normal Sevfert 1 galaxies. then similar variability should be observed in both svystenis with similar black holes.," If the X-ray variability is caused by some global coherent oscillation for both LLAGNs and normal Seyfert 1 galaxies, then similar variability should be observed in both systems with similar black holes."694 Indeed. those LLACNs and AGNs follow the same trend in Figure 1.. though LLACGNs have," Indeed, those LLAGNs and AGNs follow the same trend in Figure \ref{fig-1}, , though LLAGNs have"695Recently. there has been much interest generated by the observation of ubiquitous propagating Alfvénnic waves in the solar corona detected by Tomcezyketal.(2007) using the innovative Coronal Multi-Channel Polarimeter (CoMDP) instrument.,"Recently, there has been much interest generated by the observation of ubiquitous propagating Alfvénnic waves in the solar corona detected by \citet{tomczetal07} using the innovative Coronal Multi-Channel Polarimeter (CoMP) instrument."696 The Alfvénnic properties of (hese waves are undeniable since thev have a phase speed of about 1 Mam Ll (he velocity components are perpendicular to the direction of magnetic field lines and," The Alfvénnic properties of these waves are undeniable since they have a phase speed of about 1 Mm $^{-1}$, the velocity components are perpendicular to the direction of magnetic field lines and"697wind density and velocity and the ionization front parameters.,wind density and velocity and the ionization front parameters.698 Their numerical calculations also assume a hemispherical proplyd head and a cylindrical tail., Their numerical calculations also assume a hemispherical proplyd head and a cylindrical tail.699" Studying the interaction between these two winds, they have successfully reproduced the arc emission for the proplyds near 6! Ori C. Richling&Yorke(2000) performed 2D, axisymmetric hydrodynamical simulations of photoevaporating disks including both ionizing (hv>13.6 eV) and dissociating (6 eV < hy< 13.6 eV) radiation."," Studying the interaction between these two winds, they have successfully reproduced the arc emission for the proplyds near $\theta^1$ Ori C. \cite{rich..00} performed 2D, axisymmetric hydrodynamical simulations of photoevaporating disks including both ionizing $\nu \geq 13.6$ eV) and dissociating (6 eV $<$ $\nu <$ 13.6 eV) radiation."700" In their models, disk structures are formed through collapse simulations of 1 and 2 Mo molecular clumps (Yorke&Bodenheimer,1999),, which are then exposed to the radiation field by switching on the external UV radiation field in the calculation."," In their models, disk structures are formed through collapse simulations of 1 and 2 $M_{\odot}$ molecular clumps \citep{yorke..99}, which are then exposed to the radiation field by switching on the external UV radiation field in the calculation."701" They studied the effects of distance from the UV photon source on emission line maps and the effect of the presence of a spherical wind from the proplyd star, which promotes the appearance of collimated, bipolar microjets."," They studied the effects of distance from the UV photon source on emission line maps and the effect of the presence of a spherical wind from the proplyd star, which promotes the appearance of collimated, bipolar microjets."702" Instead of calculating the dissociation of the H5 molecule, they followed the ionization of C. They took into account photoelectric heating and cooling by fine-structure lines such as [CII] 158 um, [OT] 63 um and [OT] 145 um. An important point in their work is the inclusion of the diffuse radiation field,which is responsible for the tails of the proplyds (see Figure 1))."," Instead of calculating the dissociation of the $_2$ molecule, they followed the ionization of C. They took into account photoelectric heating and cooling by fine-structure lines such as ] 158 $\mu$ m, ] 63 $\mu$ m and ] 145 $\mu$ m. An important point in their work is the inclusion of the diffuse radiation field,which is responsible for the tails of the proplyds (see Figure \ref{f1}) )."703 In this work we present the first fully three-dimensional numerical simulations of disks exposed to FUV and EUV radiation fields., In this work we present the first fully three-dimensional numerical simulations of disks exposed to FUV and EUV radiation fields.704" As discussed in Johnstoneetal.(1998),, the diffuse ionizing and dissociating radiation field is important for determining the shape of the ionization front behind the (directly) illuminated disk surface."," As discussed in \cite{john..98}, the diffuse ionizing and dissociating radiation field is important for determining the shape of the ionization front behind the (directly) illuminated disk surface."705" Our present simulations do not include the diffuse field, resulting in the tails that do not have the correct morphology (see Cerqueira et al."," Our present simulations do not include the diffuse field, resulting in the tails that do not have the correct morphology (see Cerqueira et al."706 2006b)., 2006b).707" We concentrate on obtaining a description of the head of the proplyd flow, and study the effect of different orientations of the disks with respect to the impinging UV photon field (and the stellar wind)."," We concentrate on obtaining a description of the head of the proplyd flow, and study the effect of different orientations of the disks with respect to the impinging UV photon field (and the stellar wind)."708 The present paper is organized as follows., The present paper is organized as follows.709" In §2,, we describe the code and the simulations."," In \ref{simulations}, we describe the code and the simulations."710" In §3,, we present our results."," In \ref{results}, we present our results."711" Finally, in $4,, we draw our main conclusions."," Finally, in \ref{conclusions}, we draw our main conclusions."712" The 3D numerical simulations have been carried out with the YGUAZÜ-A adaptive grid code (Ragaetal,2000,2002) using a 5-level binary adaptive grid."," The 3D numerical simulations have been carried out with the YGUAZÚ–A adaptive grid code \citep{raga..00,raga..02} using a 5-level binary adaptive grid."713 The YGUAZU- code integrates the gasdynamic equations employing the flux vector splitting scheme of vanLeer(1982) together with a system of rate equations for atomic/ionic species.," The YGUAZ\'U--A code integrates the gasdynamic equations employing the flux vector splitting scheme of \cite{vanleer82}714 together with a system of rate equations for atomic/ionic species."715" In our simulations, we consider 4 species:HI, HII, and CII."," In our simulations, we consider 4 species:, , and ."716" This code has been extensively employed for simulating different astrophysical flows such as jets (Masciadrietal.,2002;Cerqueiraetal., 2006a),, interacting winds (Gonzálezetal., 2004),, photoevaporating clumps (Cerqueiraetal.,2006b) and supernova remnants (Velázquezetal.,20012,2004)."," This code has been extensively employed for simulating different astrophysical flows such as jets \citep{masciadri..02, dri..06a}, interacting winds \citep{ricardo..04}, photoevaporating clumps \citep{dri..06b} and supernova remnants \citep{pablo..01a,pablo..04}."717". It was also tested with laser generated plasma laboratory experiments (Sobraletal.,2000;Raga2001;Velázquez2001b)."," It was also tested with laser generated plasma laboratory experiments \citep{sobral..00,raga..01,pablo..01b}."718". In this work, instead of solving an energy equation, we prescribe a temperature law, given by where T,=10000 K, is the characteristic temperature of an region, T?=1000 K, the typical PDR temperature, K, the temperature of the molecular gas, xy is the hydrogen ionization fraction and xcy is the carbon ionization fraction (xcg=1 when all the CI is CID)."," In this work, instead of solving an energy equation, we prescribe a temperature law, given by where $T_1 = 10\,000$ K, is the characteristic temperature of an region, $T_2 = 1\,000$ K, the typical PDR temperature, $T_3 = 10$ K, the temperature of the molecular gas, $x_\mathrm{H II}$ is the hydrogen ionization fraction and $x_\mathrm{C719II}$ is the carbon ionization fraction $x_\mathrm{C II}=1$ when all the $\mathrm{C I}$ is $\mathrm{C II}$ )."720" This prescription is justified if the thermal equilibrium time scale is much smaller than the dynamical time scale (Lefloch&Lazareff,1994),, which is the case here, for both the ionized and the molecular gas."," This prescription is justified if the thermal equilibrium time scale is much smaller than the dynamical time scale \citep{lefloch..94}, which is the case here, for both the ionized and the molecular gas."721" It is clear from equation (1)), that we have possible temperatures ranging from 10 K to 10 K. In order to study the dependence of the PDR geometry with temperature, we also compute models in which we set 7;=3000 K. Following Richling&Yorke(2000) we do not treat thedissociation of Ho, but consider that the dissociation front and the carbonionization front coincide."," It is clear from equation \ref{temp}) ), that we have possible temperatures ranging from 10 K to $10^4$ K. In order to study the dependence of the PDR geometry with temperature, we also compute models in which we set $T_2 = 3\,000$ K. Following \cite{rich..00} we do not treat thedissociation of $_2$ but consider that the dissociation front and the carbonionization front coincide."722 We then solve rate equations for hydrogen and carbon takinginto, We then solve rate equations for hydrogen and carbon takinginto723"from the kinetic energv in a buller region nine cells (2935.tkpe) thick outside of ri; where the stun is over the jV; cells in the buffer region rj,<ry«rv,+9f.",from the kinetic energy in a buffer region nine cells $h^{-1}$ kpc) thick outside of $r_{vir}$: where the sum is over the $N_{b}$ cells in the buffer region $r_{vir}<r_k<r_{vir}+9l$.724 Now suppose the gas is allowed to rearrange itself within the DM potential such that ibis in hydrostatic equilibrium and has a polvtropic equation of state with index P—1.2., Now suppose the gas is allowed to rearrange itself within the DM potential such that it is in hydrostatic equilibrium and has a polytropic equation of state with index $\Gamma=1.2$.725 There is much support for such a model— see (he discussion and Fie., There is much support for such a model— see the discussion and Fig.726 1 in Drvan).. and also Ascasibaretal.(2006).," 1 in \citet[][comparing the polytropic model with a full, high resolution 727cosmological simulation by G. Bryan]{OstrikerBB05}, and also \citet{AscasibarSYMG06}."728. We will treat the gas as a tracer such that the potential. set bv the DM. does not change.," We will treat the gas as a tracer such that the potential, set by the DM, does not change."729" Defining as in Ostrikerοἱal.(2005).. the resulting gas pressure 7 and density p are given by llere 0,4 is a nonthermal component of pressure. assumed to be proportional to thermal pressure such that the total 2,=(12-9,,)P."," Defining as in \citet{OstrikerBB05}, the resulting gas pressure $P$ and density $\rho$ are given by Here $\delta_{rel}$ is a nonthermal component of pressure, assumed to be proportional to thermal pressure such that the total $P_{tot}=(1+\delta_{rel})P$."730" To specify the final gas distribution given (hese assumptions. two quanliGes still need to be determined. namely the pressure 71, and densitv py at the potential minimum."," To specify the final gas distribution given these assumptions, two quantities still need to be determined, namely the pressure $P_0$ and density $\rho_0$ at the potential minimum."731 This can be done with (wo equations of constraint. derived bv requiring conservation of energv and by matching the external surface pressure. as follows.," This can be done with two equations of constraint, derived by requiring conservation of energy and by matching the external surface pressure, as follows."732"For a given choice of [η and po. the final radius ry of the gas initially inside μις CAL be found by summing outwards [rom the cluster center until the initial mass M, is enclosed: This implies that eas mav expand or contract. changing (he gas [fraction inside rj.","For a given choice of $P_0$ and $\rho_0$, the final radius $r_f$ of the gas initially inside $r_{vir}$ can be found by summing outwards from the cluster center until the initial mass $M_g$ is enclosed: This implies that gas may expand or contract, changing the gas fraction inside $r_{vir}$."733 Assuming the external surface pressure changes little with radius. there will be mechanical work done. causing a change in energy proportional to the change in volume. M. The equation lor conservation of energy is (hus," Assuming the external surface pressure changes little with radius, there will be mechanical work done, causing a change in energy proportional to the change in volume, $\Delta E_p=(4\pi/3)(r_{vir}^3-r_f^3)P_s$ The equation for conservation of energy is thus"734determined by the current conditions in the star-disk interaction and independent of any initial conditions.,determined by the current conditions in the star-disk interaction and independent of any initial conditions.735" MP05 described the method of computing the stellar spin rate in this hypothetical equilibrium state for their torque formulation, which uses equations (23)-(27) of that work."," MP05 described the method of computing the stellar spin rate in this hypothetical equilibrium state for their torque formulation, which uses equations (23)–(27) of that work."736" In presenting our results below, we compare the actual spin rates to the equilibrium values to gain insight into these systems."," In presenting our results below, we compare the actual spin rates to the equilibrium values to gain insight into these systems."737" The coupled equations (1)), (2)), and (3)) describe the evolution of the system."," The coupled equations \ref{eq_mdot}) ), \ref{eq_rstar}) ), and \ref{eq_angmom}) ) describe the evolution of the system."738" We wrote a computational code that solves these simultaneously, using the fourth-order Runge-Kutta scheme of Pressetal.(1994),, starting from ty=3x104 yr and ending at 3 Myr."," We wrote a computational code that solves these simultaneously, using the fourth-order Runge-Kutta scheme of \citet{pressea94}, starting from $t_0=3\times10^4$ yr and ending at 3 Myr."739" For computational efficiency and numerical stability, we use a dynamic timestep."," For computational efficiency and numerical stability, we use a dynamic timestep."740" At each step, we compute the next timestep by requiring that none of the three main variables, Ὡς, R,, M,, change by more than per step."," At each step, we compute the next timestep by requiring that none of the three main variables, $\Omega_{*}$, $R_{*}$, $M_{*}$, change by more than per step."741 This gives results converged to three significant figures and typically requires a few hundred timesteps for the entire evolution., This gives results converged to three significant figures and typically requires a few hundred timesteps for the entire evolution.742" In order to compute the torque at each timestep, the code first solves equation to determine the magnetic connection state of the system."," In order to compute the torque at each timestep, the code first solves equation \ref{eq_fcrit}) ) to determine the magnetic connection state of the system."743"(8)) Then, depending on the state, the code solves either equation (10)) or (11)) to determine the disk truncation radius."," Then, depending on the state, the code solves either equation \ref{eq_rt1}) ) or \ref{eq_rt2}) ) to determine the disk truncation radius."744" We enforce a minimum value of R;=R,.", We enforce a minimum value of $R_t=R_*$.745" Finally, the code calculates the accretion torque using equation (12)) and the magnetic torque using either equation (13)) or "," Finally, the code calculates the accretion torque using equation \ref{eq_ta}) ) and the magnetic torque using either equation \ref{eq_tmag2}) ) or \ref{eq_tmag1}) )."746The torque is assumed to be constant during each timestep., The torque is assumed to be constant during each timestep.747"(14)). We also enforce a maximum on the spin rate, corresponding to f=1 (see §2.3))."," We also enforce a maximum on the spin rate, corresponding to $f=1$ (see \ref{sec_spin}) )."748 This section contains results from several models., This section contains results from several models.749" Table 1 lists the parameters for all cases, in order of their presentation and grouped by the figures in which the results appear."," Table \ref{tab_parms} lists the parameters for all cases, in order of their presentation and grouped by the figures in which the results appear."750" All models have the same initial stellar radius and effective temperature (Τε=4280 K; see (8E)reffig, starand§ 2.2)).", All models have the same initial stellar radius $8R_\odot$ ) and effective temperature $T_e=4280$ K; see \\ref{fig_rstar} and \ref{sec_rstar}) ).751" Also, all models end with a mass of Mo, so that the initial mass depends on the accretion rate (see 2.1))."," Also, all models end with a mass of $M_\odot$, so that the initial mass depends on the accretion rate (see \ref{sec_mdot}) )."752" For each case in table 1,, we ran four models, one for each combination of two different initial spin rates (see 82.3)) and two different mass accretion rates (see 2.1)."," For each case in table \ref{tab_parms}, we ran four models, one for each combination of two different initial spin rates (see \ref{sec_spin}) ) and two different mass accretion rates (see \ref{sec_mdot}) )."753 These parameters are intended to represent a range that is appropriate for accreting T Tauri stars., These parameters are intended to represent a range that is appropriate for accreting T Tauri stars.754" Section 3.1 includes two simplified cases with no magnetic fields: one in which T;=0, so that there is no star-disk interaction; and one in which B,=0 so that there are no magnetic effects, just disk accretion."," Section \ref{sec_t0} includes two simplified cases with no magnetic fields: one in which $T_*=0$, so that there is no star-disk interaction; and one in which $B_*=0$ so that there are no magnetic effects, just disk accretion."755" Section 3.2 contains the main results, models that include magnetic fields with realistic field line opening = 1) and different values of 8, for two different field strengths."," Section \ref{sec_gam1} contains the main results, models that include magnetic fields with realistic field line opening $\gamma_c=1$ ) and different values of $\beta$, for two different field strengths."756"(ye For comparison, and as a verification of our model and computations, the appendix contains results for cases with the classical assumption of a fully closed magnetic field (9.= oo) and 8—1."," For comparison, and as a verification of our model and computations, the appendix contains results for cases with the classical assumption of a fully closed magnetic field $\gamma_c=\infty$ ) and $\beta=1$."757" 'To begin, it is useful to examine few simplified cases."," To begin, it is useful to examine a few simplified cases."758" The main goal is to understand aindividually how the stellar contraction and accretion alone affects the spin evolution, before adding the effects of magnetic fields."," The main goal is to understand individually how the stellar contraction and accretion alone affects the spin evolution, before adding the effects of magnetic fields."759" We will consider two cases here: one in which the net torque on the star is forced to be zero; and one in which there are no magnetic fields, so that the only torque on the star is due to the accretion of material from the disk."," We will consider two cases here: one in which the net torque on the star is forced to be zero; and one in which there are no magnetic fields, so that the only torque on the star is due to the accretion of material from the disk."760 Figure 4 shows the evolution of stellar spin for the case of zero torque (Τε=0 in [3]])., Figure \ref{fig_t0} shows the evolution of stellar spin for the case of zero torque $T_*=0$ in ]).761 The figure shows both, The figure shows both762Dased on the IRACI colour (lower panel). we find four sources. 771. 660. and SOri665 / with >30 colour excesses. which we consider to be primary disk In all [our cases the excess increases significantly in IRAC channel 4 compared with channel 3. indicating rising flux levels towards longer wavelengths. a clear signature of disk emission.,"Based on the IRAC4 colour (lower panel), we find four sources – 71, J053949.5-023130, 60, and 65 – with $>3\sigma$ colour excesses, which we consider to be primary disk In all four cases the excess increases significantly in IRAC channel 4 compared with channel 3, indicating rising flux levels towards longer wavelengths, a clear signature of disk emission."763 With the exception of SO011665. these objects have been published previously as disk-bearing verv low mass sources (Caballeroetal. 2007)..," With the exception of 65, these objects have been published previously as disk-bearing very low mass sources \citep{2007A&A...470..903C, 2007A&A...472L...9Z}."764 This gives a disk fraction of 4 out of 14 or 29-054..., This gives a disk fraction of 4 out of 14 or $29\pm ^{16}_{13}$.765 Here we do not count the three objects with upper limits well-above the photospheric level. for which we cannot decide if they have disk excess or not.," Here we do not count the three objects with upper limits well-above the photospheric level, for which we cannot decide if they have disk excess or not."766 Disk frequencies derived Irom IRACS and IRAC are thus consistent; since the value determined from [RACH is likely to be more robust. we pul more emphasis on this result.," Disk frequencies derived from IRAC3 and IRAC4 are thus consistent; since the value determined from IRAC4 is likely to be more robust, we put more emphasis on this result."767 We note that the disk fraction in our sample might still be somewhat higher than given here. due to the combined effects of photometric uncertainties and contaminating field objects (seethedicussioninCaballeroetal.2007)..," We note that the disk fraction in our sample might still be somewhat higher than given here, due to the combined effects of photometric uncertainties and contaminating field objects \citep[see the dicussion in][]{2007A&A...470..903C}."768 We now compare our disk detection rate will previous results for more massive objects in c OOri. as given by Hernándezetal.(2007). based on IRAC! data: for ILXeDe stars. [or intermediate-mass T. Tauri stars. for T Tauri stars. for brown dwarfs.," We now compare our disk detection rate with previous results for more massive objects in $\sigma$ Ori, as given by \citet{2007ApJ...662.1067H} based on IRAC data: for HAeBe stars, for intermediate-mass T Tauri stars, for T Tauri stars, for brown dwarfs."769 The value for brown chwarls is in agreement with the disk fraction of lderived by. Javawardhana [rom ground-based. L’-band imaging., The value for brown dwarfs is in agreement with the disk fraction of derived by \citet{2003AJ....126.1515J} from ground-based L'-band imaging.770 A higher brown dwarf disk fraction of has been published by Caballeroetal.(2007).., A higher brown dwarf disk fraction of has been published by \citet{2007A&A...470..903C}.771 For the IMPO range (objects with ALS201 μμ). we now derive a disk fraction of29%... which is compatible with the values for T Tauri stars and brown dwarls within (he lo uncertainties.," For the IMPO range (objects with $M\lesssim 20\,M_{\mathrm{Jup}}$ ), we now derive a disk fraction of, which is compatible with the values for T Tauri stars and brown dwarfs within the $\sigma$ uncertainties."772 We do not see a trend to hieher disk [requencies in the IPMO range. as claimecl by ZapateroOsorioetal.(2007).. instead the evidence points to comparable disk fractions for planetary mass objects. brown cdwarfs. and T Tawi stus. Le. over more (han (wo orders of magnitude in object mass (0.008...2M. ).," We do not see a trend to higher disk frequencies in the IPMO range, as claimed by \citet{2007A&A...472L...9Z}, instead the evidence points to comparable disk fractions for planetary mass objects, brown dwarfs, and T Tauri stars, i.e. over more than two orders of magnitude in object mass $0.008\ldots 2\,M_{\odot}$ )."773 We note that two of the objects with disk excess. 666 and 771. stand out [rom the rest of (he sample. as they show excessively strong Io. emission with equivalent widths ol ~100 and ~TOOA.. respectively (Darradov.Navaseuésοἱal.2001.2002)..," We note that two of the objects with disk excess, 66 and 71, stand out from the rest of the sample, as they show excessively strong $\alpha$ emission with equivalent widths of $\sim 100$ and $\sim 700$, respectively \citep{2001A&A...377L...9B,2002A&A...393L..85B}."774 This indicates that the presence of a dusty disks is likely accompanied by ongoing gas accretion. causing intense Ila emission. as observed in T Tauri stars and brown clwarts.," This indicates that the presence of a dusty disks is likely accompanied by ongoing gas accretion, causing intense $\alpha$ emission, as observed in T Tauri stars and brown dwarfs."775where L(r.8) is (he amplitude ofthe A effect and A is the inclination of the flux vector wilh respect to the rotational axis.,"where $L(r,\theta)$ is the amplitude ofthe $\Lambda$ effect and $\lambda$ is the inclination of the flux vector with respect to the rotational axis."776 We use for (he amplitude of the A effect the expressions where d=0.025..., We use for the amplitude of the $\Lambda$ effect the expressions where $d=0.025R_\odot$.777 À aud Ag are lree-parameters., $\lambda$ and $\Lambda_0$ are free-parameters.778" The value of / needs to be equal to or larger than 2 (o ensure regularitv near the pole. so we set /=2. The X effect does nol depend on ce,. ej or O4. meaning il is a stationary elfect."," The value of $l$ needs to be equal to or larger than 2 to ensure regularity near the pole, so we set $l=2$ The $\Lambda$ effect does not depend on $v_r$, $v_\theta$ or $\Omega_1$, meaning it is a stationary effect."779 We emphasize that the A ellect depends on stellar angular velocity. ο since (he X effect is generated by turbulence and Coriolis force.," We emphasize that the $\Lambda$ effect depends on stellar angular velocity $\Omega_0$, since the $\Lambda$ effect is generated by turbulence and Coriolis force."780 The more rapidly the star rotates. the more angular momentum the A effect can transport.," The more rapidly the star rotates, the more angular momentum the $\Lambda$ effect can transport."781 The dependence of Ay anc A on stellar angular velocity is discussed in 8??.., The dependence of $\Lambda_0$ and $\lambda$ on stellar angular velocity is discussed in \ref{variation}.782 Using the moclifiecl Las-Wendroll seheme with TVD artificial viscosity (Davis1984).. we solve Equations (1))-(5)) numerically for the northern hemisphere of the meridional plane in 0.652.«r0.934. and 0«9<7/2.," Using the modified Lax-Wendroff scheme with TVD artificial viscosity \citep{davis1984tvd}, we solve Equations \ref{continuity}) \ref{se1}) ) numerically for the northern hemisphere of the meridional plane in $0.65R_\odot < r <0.93R_\odot$ and $0 < \theta < \pi/2$."783 We use a uniform resolution of 200 points in the radial direction and 400points in the Iatitudinal direction in all of our simulations., We use a uniform resolution of $200$ points in the radial direction and $400$points in the latitudinal direction in all of our simulations.784 Each simulation run is conducted until it reaches a stationary state., Each simulation run is conducted until it reaches a stationary state.785" All the variables py. v. ry. O4 and s, are equal to zero in the initial condition."," All the variables $\rho_1$, $v_r$ , $v_\theta$ , $\Omega_1$ and $s_1$ are equal to zero in the initial condition."786At the top boundary (r=0.9322. ) we adopt stress-Iree boundary conditions for ος. eg aud O4 and set thederivative of δι to,"At the top boundary $r=0.93R_\odot$ ) we adopt stress-free boundary conditions for $v_r$ , $v_\theta$ and $\Omega_1$ and set thederivative of $s_1$ to"787"and 4 ryd, respectively, will propagate at the same (photon flux-limited) velocity for a critical spectral index, corresponding to our adopted He/H abundance, y= 0.0823.","and 4 ryd, respectively, will propagate at the same (photon flux-limited) velocity for a critical spectral index, corresponding to our adopted He/H abundance, $y = 0.0823$ ."788" The agreement with the HS'T--observed mean spectral index of AGN at 1-2 ryd, (as)=1.76+0.12 22002),(Telfer suggests that aand iionization fronts normally propagate together."," The agreement with the -observed mean spectral index of AGN at 1–2 ryd, $\langle \alpha_s \rangle = 1.76 \pm 0.12$ (Telfer 2002), suggests that and ionization fronts normally propagate together."789" In photoionization equilibrium in regions of high ionization, zgry«1 and wen<1, the aabundance ratio is (Fardal 11998; Shull 22004), 'The numerical coefficient has been increased from 1.70 to 1.77, reflecting an updated value of the primordial"," In photoionization equilibrium in regions of high ionization, $x_{\rm HI} \ll 1$ and $x_{\rm HeII} \ll 1$, the abundance ratio is (Fardal 1998; Shull 2004), The numerical coefficient has been increased from 1.70 to 1.77, reflecting an updated value of the primordial"790the neutron star interiors. (,the neutron star interiors. (791This is not to be confused with the global length-scale of neutron stars (o10Η) for which M/R~0.3 depending on the stars mass (in uniisec— G=1 so that M.71.475/).),This is not to be confused with the global length-scale of neutron stars $\sim 10km$ ) for which $M/R\sim 0.3$ depending on the star's mass (in units $c=G=1$ so that $M_{\sun}\approx 1.475km$ ).)792 In other words. eravily curves space-time only on a macroscopic scale but to a verv good approximation leaves it [lat on a microscopic scale.," In other words, gravity curves space-time only on a macroscopic scale but to a very good approximation leaves it flat on a microscopic scale."793 To achieve an appreciable curvature on a microscopic level at which the strong interactions dominate the particle dvnanmics mass densities greater than ~107g10em? would be necessary (WeberThorne 1966)..," To achieve an appreciable curvature on a microscopic level at which the strong interactions dominate the particle dynamics mass densities greater than $\sim 10^{40} g\hspace{1mm}cm^{-3}$ would be necessary \citep{Weber:1999a,Thorne1966a}."794 Under this circumstances the problem of constructing models of neutron stars separates into (wo distinct tasks., Under this circumstances the problem of constructing models of neutron stars separates into two distinct tasks.795 First. the short-range effects of the nuclear forces are described by (he principles of many-body nuclear physics in a local inerGal Irae. proper relerence frame) in which space-time is flat.," First, the short-range effects of the nuclear forces are described by the principles of many-body nuclear physics in a local inertial frame (co-moving proper reference frame) in which space-time is flat."796 Second. the coupling between the long-range gravitational [orce and matter is accounted for by solving the general relativistic equations for the gravitational field described by the curvature of space-time. leading to the global structure of stellar configurations.," Second, the coupling between the long-range gravitational force and matter is accounted for by solving the general relativistic equations for the gravitational field described by the curvature of space-time, leading to the global structure of stellar configurations."797 In (he case of spherically svinnetric static (non-rotating) stars the metric has (he famous Schwarzschild form: with For a static star Einsteins field equations (eq. (1))), In the case of spherically symmetric static (non-rotating) stars the metric has the famous Schwarzschild form: $(c=G=1)$ where the metric functions $\phi(r)$ and $\Lambda(r)$ are given by: with For a static star Einstein's field equations (Eq. \ref{eq.1}) ))798 reduce then to the familiar Tolman- equation (TOV) (Tolman1939:Oppenheimer&Volkolf 1939):," reduce then to the familiar Tolman-Oppenheimer-Volkoff equation (TOV) \citep{Tolman:1939jz,PhysRev.55.374}: :"799Suppose that the true redshifts 2 are available for a subset of the objects: for now. assume that the subset is a random subsample of he objects in a magnitude limited. catalog.,"Suppose that the true redshifts $z$ are available for a subset of the objects; for now, assume that the subset is a random subsample of the objects in a magnitude limited catalog."800 Ideally. this stibset would have the same geometry as the full survey. as Cross-correlating the objects with spectra and those witrout allows the use of other methods. (e.g. Calor ο al.," Ideally, this subset would have the same geometry as the full survey, as cross-correlating the objects with spectra and those without allows the use of other methods (e.g. Caler et al."801 2009)., 2009).802 In practice. this may be clillicult to achieve and this is not required. for the analysis which follows. provied that the photometric redshift estimator does not have spatially dependent. biases (ee. as a result of photometric calibrations varving across the survey).," In practice, this may be difficult to achieve – and this is not required for the analysis which follows, provided that the photometric redshift estimator does not have spatially dependent biases (e.g., as a result of photometric calibrations varying across the survey)."803 lor the οjects with spectroscopic redshifts. one can study 1ο joint distribution of ς and : (see Figure 1)).," For the objects with spectroscopic redshifts, one can study the joint distribution of $\zeta$ and $z$ (see Figure \ref{pzzeta}) )."804 Typically. most photometric redshift’ codes are constructed to return. ic]zzz2.," Typically, most photometric redshift codes are constructed to return $\langle\zeta |z\rangle \approx z$."805 The codes which do so are sometimes said to be unluased. but they are not. perfect: the scatter around the unbiased. mean is of order ay.στ0.05(1|z).," The codes which do so are sometimes said to be unbiased, but they are not perfect: the scatter around the unbiased mean is of order $\sigma_{\zeta|z} \approx 0.05\,(1+z)$."806 This scatter. ¢ombinedwith the fact that £C[z)zx2 means that £z]4¢: the fact that 2 isguaranteed to be biased is not widely appreciated.," This scatter, combinedwith the fact that $\langle\zeta |z\rangle \approx z$ means that $\langle z|\zeta\rangle \ne \zeta$: the fact that $\langle z|\zeta\rangle$ is to be biased is not widely appreciated."807 However. we show below that it matters Little whether ([z) or (z|O) are unbiased what matters is that the bias is accurately quantified.," However, we show below that it matters little whether $\langle\zeta |z\rangle$ or $\langle z|\zeta\rangle$ are unbiased – what matters is that the bias is accurately quantified."808 In particular. af ANde and ANας denote the clistribution of C and z values in the subset of he data where both z and care available. then what matters is that p(C|z) and p(z]C). where ave known.," In particular, if ${\rm d}{\cal N}/{\rm d}\zeta$ and ${\rm d}N/{\rm d}z$ denote the distribution of $\zeta$ and $z$ values in the subset of the data where both $z$ and $\zeta$ are available, then what matters is that $p(\zeta|z)$ and $p(z|\zeta)$, where are known."809 Note that The algorithm in Sheth (2007) assumes that p(c|z). measured in the subset for which both z and ¢ are available. also applies to the full sample for which z is not available.," Note that The algorithm in Sheth (2007) assumes that $p(\zeta|z)$, measured in the subset for which both $z$ and $\zeta$ are available, also applies to the full sample for which $z$ is not available."810 Since dAfede is measured in the full dataset. ancl p(c[z) is known. a deconvolution is then. used to estimate the true cLNας.," Since ${\rm d}{\cal N}/{\rm d}\zeta$ is measured in the full dataset, and $p(\zeta|z)$ is known, a deconvolution is then used to estimate the true ${\rm d}N/{\rm d}z$."811 Suppose. however. that one measured. p(z|C). instead.," Suppose, however, that one measured $p(z|\zeta)$ instead."812 Then. because one could estimate the quantity on the left hand. side by cconvolving the two measurables on the right hand side.," Then, because one could estimate the quantity on the left hand side by `convolving' the two measurables on the right hand side."813 Lor the data-subset in which both z and & are available. this is correct by definition.," For the data-subset in which both $z$ and $\zeta$ are available, this is correct by definition."814 Clearly. to use this method on the larger dataset for which only ¢ is available. one must assume that p(z|) in the subset rom which it was measured remains accurate in the larger dataset.," Clearly, to use this method on the larger dataset for which only $\zeta$ is available, one must assume that $p(z|\zeta)$ in the subset from which it was measured remains accurate in the larger dataset."815 Rossi et al. (, Rossi et al. (8162010) ave shown that the deconvolution method accurately reconstructs the true Ας distribution [rom LAY dé.,2010) have shown that the deconvolution method accurately reconstructs the true ${\rm d}N/{\rm d}z$ distribution from ${\rm d}{\cal N}/{\rm d}\zeta$ .817 Figure 2 shows that the convolution approach also works well. even when only a random of the full dataset is used to calibrate p(z|) as displaved. in," Figure \ref{Nzconv} shows that the convolution approach also works well, even when only a random of the full dataset is used to calibrate $p(z|\zeta)$ – as displayed in"818moclels in disagreement witLa distauce of dzzLs6pc.,models in disagreement with a distance of $d\approx 186$ pc.819 Án important point is he sigtificance of the two-temperature WD it over the siugle WD fit., An important point is the significance of the two-temperature WD fit over the single WD fit.820 A change in chi-squared from 1.55 o 1.38 (for theSE fits) aid. even nore so. [roi13.73 t0 3.56 (for the+IUE fits) is al bes a modest improvement iu he fi quaity with tle adcditiou of he second temperature coriponen.," A change in chi-squared from 1.58 to 1.38 (for the fits) and, even more so, from 3.73 to 3.56 (for the fits) is at best a modest improvement in the fit quality with the addition of the second temperature component."821 However. addiug the second coiiporent results in a better fit o the bottom « ‘the Ly liue aroud 1025À.. aud improves he fit to the left. winees of Ly? (the Ἱθη wing is contamüuatecd N the OVI etjlssiou feature).," However, adding the second component results in a better fit to the bottom of the $\beta$ line around 1025, and improves the fit to the left wing of $\beta $ (the right wing is contaminated by the OVI emission feature)."822 Iu that region le lmprovelrent of the fit is the actual ing of the blue wit go“tl e Lys iu the spectruil., In that region the improvement of the fit is the actual fitting of the blue wing of the $\beta$ in the spectrum.823 We remark here that there coid some instrumeut backeroutd contamination contributiug o the flux. s that the Ly wotld :willy never go to zero.," We remark here that there could be some instrument background contamination contributing to the flux, such that the $\beta$ would actually never go to zero."824" However. since we have discarded all the noisy [ions of the channels {Islally the edges). the actual coutri»utiou of the instruiient containinatio 1ould be less than z5x1015 tem 7A! («505€ of the Πιν which is the excess enmisslon a (this is an ove""esimate. since the region A«010A. is near the edge where the noise is maximal)."," However, since we have discarded all the noisy portions of the channels (usually the edges), the actual contribution of the instrument contamination should be less than $\approx 5 \times 10^{-15}$ $~$ $^{-1}$ $^{-2}$ $^{-1}$ $<$ of the flux) which is the excess emission at (this is an overestimate, since the region $\lambda < 910$ is near the edge where the noise is maximal)."825 Iu the wavelenehi rauge A«9 he iiuiprovement ofthe fit does no fit any actual feature but only reduces the discrepancy between he model aud the oervation., In the wavelength range $\lambda <$ the improvement of the fit does not fit any actual feature but only reduces the discrepancy between the model and the observation.826 Therefore the need aud importance of the secoud component does ot originate [rom fittiug that pa1 of the spectrum where there is emission but rather it comes rou [ittiug a feature in the continuum., Therefore the need and importance of the second component does not originate from fitting that part of the spectrum where there is emission but rather it comes from fitting a feature in the continuum.827 At this stage the “belt” is really a flat continuum added to improve the fit and i should be regarded as a featureless blue spectruu., At this stage the “belt” is really a flat continuum added to improve the fit and it should be regarded as a featureless blue spectrum.828 So far the belt is probably ot the best physica description of the data but it is the best availae model compoient. to help improve e fit., So far the belt is probably not the best physical description of the data but it is the best available model component to help improve the fit.829 La fact. the two-temperature WD fit does not provide the lowest X7.> tle disk+WD oes.," In fact, the two-temperature WD fit does not provide the lowest $\chi^2_{\nu}$, the disk+WD does."830 However. this lowest disk+WD model. while fittiug better ii theIUE (lower resolution) spectral ange. does uot provide a good fit ain theFUSE (higher resoution) range of the combined spectrum al is inconsistant with the clistauce of the system.," However, this lowest disk+WD model, while fitting better in the (lower resolution) spectral range, does not provide a good fit a in the (higher resolution) range of the combined spectrum and is inconsistant with the distance of the system."831 As we stated previously. we do uot chose bliidly the lowest V7 model. but we chose one of the lowest \7 uodels that provides a better Π to so110 Ἡyecilic parts aud features of the spectruu.," As we stated previously, we do not chose blindly the lowest $\chi^2_{\nu}$ model, but we chose one of the lowest $\chi^2_{\nu}$ models that provides a better fit to some specific parts and features of the spectrum."832 The [act that both the WD-belt. aud he WD-+disk provide the lowest. V7 models. reflects he fact that the secoud Component cannot. oesently. be uodeled accurately.," The fact that both the WD+belt and the WD+disk provide the lowest $\chi^2_{\nu}$ models, reflects the fact that the second component cannot, presently, be modeled accurately."833 Trere are broad. emissior lines which are probaby due to a ho eas. aud theOVI recd-shiltect eature inay luidicate the possibility that the maerial is llowing away [rom the observer.," There are broad emission lines which are probably due to a hot gas, and the red-shifted feature may indicate the possibility that the material is flowing away from the observer."834 However. he other eiisslou features a‘e not resolved. enctela o confirm or 'efute such a scenario.," However, the other emission features are not resolved enough to confirm or refute such a scenario."835 We cdo 100 discuss here the origin or he possible scenariOs of such a flow. tlough the mechanisms at work could ye as varied as the ones discussed in Hoa«|οἱal.(2003) [9]1 the FUYV observation of the complex system DW UMa.," We do not discuss here the origin or the possible scenarios of such a flow, though the mechanisms at work could be as varied as the ones discussed in \citet{hoa03} on the FUV observation of the complex system DW UMa."836 I melt be worth iotiug that the supra-solar Nitrogen abundauce. albeit uncertain (as it. could well be from intersellar origin). togetler with the sub-solar Carbon abundance. could be a result of CNO-processiug. either from a past nova or from CNO processed," It might be worth noting that the supra-solar Nitrogen abundance, albeit uncertain (as it could well be from interstellar origin), together with the sub-solar Carbon abundance, could be a result of CNO-processing, either from a past nova or from CNO processed"837D) js dsed that encodes replicated observations for hi calilxator and tarect.,(OB) is used that encodes replicated observations for both calibrator and target.838 Raod star acquisition Wwuch is iuportaut for ensuring eood calibration — cau be angleited by usine a simple offset o* the telescope without liaving to preset the telescope and perform a ful (iustameutal) re-acquisitiou., Rapid star acquisition -- which is important for ensuring good calibration – can be augmented by using a simple offset of the telescope without having to preset the telescope and perform a full (instrumental) re-acquisition.839" In this 110do. which las been chisched ""star hopping”. t1ο adaptive optics loop is opened. while the template orders a dither to bring a different sar into the AO field selector."," In this mode, which has been christened “star hopping”, the adaptive optics loop is opened, while the template orders a dither to bring a different star into the AO field selector."840 The AO is then closed iuamally by the operator without incurring the tine peatv for re-optimization., The AO is then closed manually by the operator without incurring the time penalty for re-optimization.841 Star hopping therefore onv works ou objects of conrpirable brightucss in the wavetroit SCLISOL., Star hopping therefore only works on objects of comparable brightness in the wavefront sensor.842 As t1e observation progresses. the repetition of the template collects cight datacubes of multiple frames (typically a hundred). cach at a differcut dither position on the ¢ctector.," As the observation progresses, the repetition of the template collects eight datacubes of multiple frames (typically a hundred), each at a different dither position on the detector."843 For these observations. the detector was windowed ο 512«511 pixels.," For these observations, the detector was windowed to $512\times 514$ pixels."844 A siele snapshot observation vields Fourier coverage of 21 sdalial frequencies., A single snapshot observation yields Fourier coverage of 21 spatial frequencies.845 IHowewer. the pupiltrackiug mode POSIts in sky rotation on the detector as the parallacic anele changes.," However, the pupil-tracking mode results in sky rotation on the detector as the parallactic angle changes."846 This variation with time results sweeps the baselines iuto circular Fourier tracks aud permits roational aperture svuthesis tecliniques to assist with the fllius of the spatial frequency UV-plane plane. as illustrated in Fig. 2..," This variation with time results sweeps the baselines into circular Fourier tracks and permits rotational aperture synthesis techniques to assist with the filling of the spatial frequency UV-plane plane, as illustrated in Fig. \ref{uv}."847 To the authors knowledge. two ¢ata reduction software libraries-. are preseutly oreuce SAM daa.," To the author's knowledge, two data reduction software libraries are presently to reduce SAM data."848 One pipeline las been developed w Svney University. Cornell University. and Caltech from 20H ouwards. based on an earlier pix‘line from Berkeley.," One pipeline has been developed by Sydney University, Cornell University, and Caltech from 2004 onwards, based on an earlier pipeline from Berkeley."849 It iis already beenused for several pa])ers arising predomunatly youn Keck aperture nasking data (7)., It has already been used for several papers arising predominantly from Keck aperture masking data \citep{2000PASP..112..555T}.850 This reduction algorithin is based on fast Emer traustfonu (FFT)., This reduction algorithm is based on fast Fourier transform (FFT).851 The daa presented in this Po have been reduced by the sparse aIOLTULC uode pip¢‘line (SAMD) developed at the Observatoire de Paris., The data presented in this paper have been reduced by the sparse aperture mode pipeline (SAMP) developed at the Observatoire de Paris.852 This software is similar to that Toni Svduev university. ando nunnerous tests ando cross-checks have xoduced similar results.," This software is similar to that from Sydney university, and numerous tests and cross-checks have produced similar results."853 For voth pipelines. data reduction follows a similar oath:," For both pipelines, data reduction follows a similar path:"854The first billion vears after the Big Bang represents a period of great. interest Lor studies of both galaxy. formation and he evolution of the Universe as à whole.,The first billion years after the Big Bang represents a period of great interest for studies of both galaxy formation and the evolution of the Universe as a whole.855 Εις period. sees he formation of the first galaxies (2). and. consequently. he beginning and completion of the process of reionizing he Universe (77) as a result. of the copious number of ionizing photons emitted. by these sources.," This period sees the formation of the first galaxies \citep{wise_resolving_2008} and, consequently, the beginning and completion of the process of reionizing the Universe \citep{loeb_reionization_2001,loeb_frontier_2009} as a result of the copious number of ionizing photons emitted by these sources."856" Current and uture facilities aim to probe this epoch of the Universe roth using traditional methods such as surveying [faint galaxies (e.g. the James Webb Space ""Telescope: ?2)) and using novel techniques such as 210m cosmology (7) to probe he distribution of neutral hydrogen during the process of reionization.", Current and future facilities aim to probe this epoch of the Universe both using traditional methods such as surveying faint galaxies (e.g. the James Webb Space Telescope; \citealt{gardner_james_2009}) ) and using novel techniques such as 21cm cosmology \citep{furlanetto06a} to probe the distribution of neutral hydrogen during the process of reionization.857 Understanding this epoch of the Universe from a theoretical perspective therefore requires an understanding both of the sources of ionizing photons and of the thermal and ionization state of the intergalactic medium (16M) at these times., Understanding this epoch of the Universe from a theoretical perspective therefore requires an understanding both of the sources of ionizing photons and of the thermal and ionization state of the intergalactic medium (IGM) at these times.858 Additionally. the thermal and ionization history of the IGAL as a function of cosmic redshift. z. strongly alfects the visibility of the most distant galaxies ancl quasars (??7).. and the feedback exerted on the formation of new galaxios (222727272727272727727)..," Additionally, the thermal and ionization history of the IGM as a function of cosmic redshift, $z$, strongly affects the “visibility” of the most distant galaxies and quasars \citep{madau_radiative_1995,meiksin_colour_2006,dayal_visibility_2011}, and the feedback exerted on the formation of new galaxies \citep{efstathiou_suppressing_1992,quinn_photoionization_1996,navarro_effects_1997,barkana_photoevaporation_1999,bullock_reionization_2000,somerville_can_2002,859Benson:02a,Benson:02b,koposov_quantitative_2009,munoz_probing_2009,busha_impact_2010,maccio_luminosity_2010}."860 Phe process of reionization is expected. to begin with the formation of ionized bubbles around Luminous sources in the redshift range z=10 , The process of reionization is expected to begin with the formation of ionized bubbles around luminous sources in the redshift range $z=10$ --20.861These bubbles will eventually erow in size and number, These bubbles will eventually grow in size and number862"For heterodyne interferometry, one gets: ACTIWe noisecan see that the ratio Neg/Np is ΝΕΤΗΙalways less than one, which gives a clear advantage to direct imaging from the strict point of view of the noise.","For heterodyne interferometry, one gets: We can see that the ratio $N_\mathrm{eq}/N_h$ is always less than one, which gives a clear advantage to direct imaging from the strict point of view of the noise."863" This ratio appears squared in the ratio of imaging to bolometric interferferometry and without power in the ratio of imaging to heterodyne interferometry but in the latter case, the NET ratio is also less than one, penalising heterodyne interferometry."," This ratio appears squared in the ratio of imaging to bolometric interferferometry and without power in the ratio of imaging to heterodyne interferometry but in the latter case, the NET ratio is also less than one, penalising heterodyne interferometry."864" The number of equivalent baselines for a square horn array is: If one averages over directions in the baseline plane at a given |u|, a good approximation of Neg as a function of £ is given by (see Fig. 1)):"," The number of equivalent baselines for a square horn array is: If one averages over directions in the baseline plane at a given $|\vec{u}|$, a good approximation of $N_\mathrm{eq}$ as a function of $\ell$ is given by (see Fig. \ref{uapprox}) ):"865 One finally finds that a good approximation of the sensitivity ratio is: AC;and:noise These approximate formulae have been NETurcompared with actual calculations of the number of equivalent baselines for square arrays., One finally finds that a good approximation of the sensitivity ratio is: and: These approximate formulae have been compared with actual calculations of the number of equivalent baselines for square arrays.866" We have chosen 256 horns for the comparison and we compare bolometric and heterodyne interferometers with imagers having a low angular resolution of one degree, BICEP-like (Yoonetal, 2006)) and a high one of 10 arcminutes, Clover-like (Northetal, 2008))."," We have chosen 256 horns for the comparison and we compare bolometric and heterodyne interferometers with imagers having a low angular resolution of one degree, BICEP-like \cite{bicep}) ) and a high one of 10 arcminutes, Clover-like \cite{clover}) )."867 The results are shown in Fig. 2.., The results are shown in Fig. \ref{comparison}.868" We have chosen to only consider the multipole region between 0 and 200 as for higher multipoles, interferometers are less sensitive due the loss of coherence between largely separated horns."," We have chosen to only consider the multipole region between 0 and 200 as for higher multipoles, interferometers are less sensitive due the loss of coherence between largely separated horns."869 Note that the effect of coherence loss for the long baselines and the bandwidth smearing have not been taken into account here and might have a significant effect., Note that the effect of coherence loss for the long baselines and the bandwidth smearing have not been taken into account here and might have a significant effect.870 The sensitivities of the three different techniques only differ in the way the instrument filters the multipoles observed in the sky., The sensitivities of the three different techniques only differ in the way the instrument filters the multipoles observed in the sky.871 An imager is affected by its resolution on the sky while an interferometer is affected by the ratio between the number of equivalent baselines and the number of horns as a function of multipoles., An imager is affected by its resolution on the sky while an interferometer is affected by the ratio between the number of equivalent baselines and the number of horns as a function of multipoles.872 All of these filtering factors are less than one., All of these filtering factors are less than one.873" However, imagers are usually operated in such a way that they are not limited by their angular resolution in the multipole region of interest, in that case By~1, and the imager is always more sensitive than an interferometer (bolometric or heterodyne)."," However, imagers are usually operated in such a way that they are not limited by their angular resolution in the multipole region of interest, in that case $B_\ell\simeq 1$, and the imager is always more sensitive than an interferometer (bolometric or heterodyne)."874" From the strict point of view of sensitivity, interferometers can therefore only compete with low angular resolution imagers."," From the strict point of view of sensitivity, interferometers can therefore only compete with low angular resolution imagers."875 There is a large difference in sensitivity between bolometric and heterodyne interferometers compared to an imager: the ratio Neg/Nn acts quadratically on the variance for a bolometric interferometer while it acts linearly for a heterodyne instrument., There is a large difference in sensitivity between bolometric and heterodyne interferometers compared to an imager: the ratio $N_\mathrm{eq}/N_h$ acts quadratically on the variance for a bolometric interferometer while it acts linearly for a heterodyne instrument.876" This is due to the fact that with a heterodyne interferometer, equivalent baselines are averaged after their measurement, resulting in a 1/Neq factor on the variances."," This is due to the fact that with a heterodyne interferometer, equivalent baselines are averaged after their measurement, resulting in a $1/N_\mathrm{eq}$ factor on the variances."877" In a bolometric interferometer, the signals from all N; horns are added together multiplying the noise variance by N; while the coherent summation of equivalent baselines performs an efficient 1/N2, reduction of the noise."," In a bolometric interferometer, the signals from all $N_h$ horns are added together multiplying the noise variance by $N_h$ while the coherent summation of equivalent baselines performs an efficient $1/N_\mathrm{eq}^2$ reduction of the noise."878 This finally results in a factor N;/2N.q for the variance of a bolometric interferometer relative to a heterodyne one., This finally results in a factor $N_h/2N_\mathrm{eq}$ for the variance of a bolometric interferometer relative to a heterodyne one.879 This is largely compensated by the difference in NET between bolometric instruments and coherent ones., This is largely compensated by the difference in NET between bolometric instruments and coherent ones.880" When comparing them, the ratio of their NET also appears quadratically and favours bolometric instruments that are dominated by the photon noise rather than by that of the amplifiers."," When comparing them, the ratio of their NET also appears quadratically and favours bolometric instruments that are dominated by the photon noise rather than by that of the amplifiers."881 This situation may change in the future with the improvements of the HEMT technologies but at frequencies around and above 100 GHz we are unlikely to face photon noise limited HEMTs in the near future., This situation may change in the future with the improvements of the HEMT technologies but at frequencies around and above 100 GHz we are unlikely to face photon noise limited HEMTs in the near future.882 The difference between the NET would be even greater in space where the bolometers NET would drop as the background temperature while that of the coherent instruments would remain roughly constant., The difference between the NET would be even greater in space where the bolometers NET would drop as the background temperature while that of the coherent instruments would remain roughly constant.883" With the present technologies of bolometers and coherent amplifiers, the hierarchy in terms of sensitivity between the three techniques (and layout) studied here is very clear for the multipole range 25«€200 where the primordial B-mode signal is expected to be maximal."," With the present technologies of bolometers and coherent amplifiers, the hierarchy in terms of sensitivity between the three techniques (and layout) studied here is very clear for the multipole range $25<\ell<200$ where the primordial B-mode signal is expected to be maximal."884" Imagers are the most sensitive, bolometric interferometers have a lower sensitivity, the ratio dropping quadratically with the multipole considered."," Imagers are the most sensitive, bolometric interferometers have a lower sensitivity, the ratio dropping quadratically with the multipole considered."885 Heterodyne interferometers have an even lower sensitivity but the ratio with an imager drops less rapidly., Heterodyne interferometers have an even lower sensitivity but the ratio with an imager drops less rapidly.886" They remain however less sensitive than bolometric interferometers in the range of multipoles considered here, where the largest primordial B-mode signal is expected and where the lensing"," They remain however less sensitive than bolometric interferometers in the range of multipoles considered here, where the largest primordial B-mode signal is expected and where the lensing"887Cluster major mergers are among5 the most energetic5 phenomena in the Universe.,Cluster major mergers are among the most energetic phenomena in the Universe.888"ὃν They release a total energy. of the order of 10510"".1 erg. and it Mis nowadays accepte that they are the key ingredient to explain the origin and rarity of radio halos in galaxy clusters: shocks anc turbulence are generated during such energetic events. anc they deeply alfect the thermal and nonthermal properties of the intracluster medium (ICM)."," They release a total energy of the order of $10^{63}-10^{64}$ erg, and it is nowadays accepted that they are the key ingredient to explain the origin and rarity of radio halos in galaxy clusters: shocks and turbulence are generated during such energetic events, and they deeply affect the thermal and non–thermal properties of the intracluster medium (ICM)."889 Radio halos are the signposts of the nontherma components in galaxy clusters., Radio halos are the signposts of the non–thermal components in galaxy clusters.890 They are diffuse racio sources. whose size and morphology are similar to those of the underlving hot ICM (e.g. Ferrari et al.," They are diffuse radio sources, whose size and morphology are similar to those of the underlying hot ICM (e.g. Ferrari et al."891 2008. Cassano 2009 ancl Venturi 2011 for recent reviews).," 2008, Cassano 2009 and Venturi 2011 for recent reviews)."892" Their spectrum (defined. as Sxf£ "") is steep. with typical values of the spectral index o in the range 1.2.1.4."," Their spectrum (defined as $\propto\nu^{-\alpha}$ ) is steep, with typical values of the spectral index $\alpha$ in the range 1.2–1.4."893 However. recent highsensitivity low frequency. imaging led to the discovery. of racio halos with much steeper spectra (Venturi 2011). with spectral index à~LS2 (e.g. 5521. Brunetti et al.," However, recent high--sensitivity low frequency imaging led to the discovery of radio halos with much steeper spectra (Venturi 2011), with spectral index $\alpha \sim 1.8-2$ (e.g. 521, Brunetti et al."894 2008. Dallacasa ct al.," 2008, Dallacasa et al."895 2000: 66907 Macario ct al., 2009; 697 Macario et al.896 2010)., 2010).897 Combined radio anc Xrav studies. provide strong support to the idea that radio halos are found only in unrelaxed clusters., Combined radio and X–ray studies provide strong support to the idea that radio halos are found only in unrelaxed clusters.898" Buote (2001). first showed a correlation between the 1.4 Gllz radio power of halos. Piocu. and the dipole power ratio P, /Po in the hosting cluster: based on temperature maps. Govoni ct al."," Buote \cite{buote01} first showed a correlation between the 1.4 GHz radio power of halos, $_{\rm 1.4~GHz}$, and the dipole power ratio $_{\rm 1}$ $_{\rm 0}$ in the hosting cluster; based on temperature maps, Govoni et al."899 (2004). found evidence for merging activity in clusters with radio halos.," \cite{govoni04}900 found evidence for merging activity in clusters with radio halos."901 Venturi et al. (, Venturi et al. (9022008. hereinafter. VOS) showed that all racio halos in the GMIEE (Giant. Metrewave Radio Telescope) radio halo survey are located in clusters with signs of dynamical clisturbances.,"2008, hereinafter V08) showed that all radio halos in the GMRT (Giant Metrewave Radio Telescope) radio halo survey are located in clusters with signs of dynamical disturbances."903 More recently. Cassano ct al. (," More recently, Cassano et al. ("9042010. hereinafter. C10) carried out a quantitative analysis of the radio halocluster merger scenario.,"2010, hereinafter C10) carried out a quantitative analysis of the radio halo–cluster merger scenario."905 μον. used all clusters in. the GMBIRE radio halo cluster sample (Venturi et al., They used all clusters in the GMRT radio halo cluster sample (Venturi et al.906 2007. hereinafter VOT. and. VOS) with available high quality images (a total of 32. clusters) to characterize the presence of substructures by three cillerent methods.," 2007, hereinafter V07, and V08) with available high quality images (a total of 32 clusters) to characterize the presence of substructures by three different methods."907 Γον showed that clusters with and without radio halos are well segregated according to all parameters indicating substructure: raclio halos are associated with clusters currently undergoing a merger. while clusters without radio halo are usually more “relaxed”.," They showed that clusters with and without radio halos are well segregated according to all parameters indicating substructure: radio halos are associated with clusters currently undergoing a merger, while clusters without radio halo are usually more “relaxed”."908 Four clusters. however. are noticeable outliers in the correlations. being disturbed systems with no detectable radio halo at the sensitivity limit of the 610. MlIIz GAIRT survey (VOT and VOS).," Four clusters, however, are noticeable outliers in the correlations, being disturbed systems with no detectable radio halo at the sensitivity limit of the 610 MHz GMRT survey (V07 and V08)."909 One o£ the outliers. Abell 781," One of the outliers, Abell 781"910heir coalescence driven by the emission of eravitational waves. and the recoil associated with the non-zero uct Inear momentum carried away by CAVs in the coalescence of two unequal mass black holes (the veravitational rocket,"their coalescence driven by the emission of gravitational waves, and the recoil associated with the non-zero net linear momentum carried away by GWs in the coalescence of two unequal mass black holes (the “gravitational rocket”)."911 Major halo mergers lead to MDII fueling aud rigecr ).quasar activity., Major halo mergers lead to MBH fueling and trigger quasar activity.912 In this paper we use the same uodel to provide a more detailed characterization of the CV. sienal from iuspiraliug MDIIDs., In this paper we use the same model to provide a more detailed characterization of the GW signal from inspiraling MBHBs.913 Their coutributiou o the data stream is twofold: unresolved. sources will eive origin to confusion noise to be compared to instrumental noise aud other astrophysical stochastic vackerounds (e.g. from white dwarf binaries. Financer Phinney 2003). while resolved iuspialiug binaries will xobe eravitv in extreme conditions (e.e.. Vecchio 2001).," Their contribution to the data stream is twofold: unresolved sources will give origin to confusion noise to be compared to instrumental noise and other astrophysical stochastic backgrounds (e.g. from white dwarf binaries, Farmer Phinney 2003), while resolved inspiraling binaries will probe gravity in extreme conditions (e.g., Vecchio 2004)."914 Confusion noise aud resolved sources should provide different cosmological information., Confusion noise and resolved sources should provide different cosmological information.915 The former. produced wea large nuniber of unresolved MDIIDs. will trace light AIBIIDs at very high redshift. placing coustraiuts ou black tole formation scenarios prior to the reionization epoch: he latter will be a formidable tool to follow the cosmic evolution of MDIIs aud the formation and dvnamics of MBI binaries following galaxy mergers.," The former, produced by a large number of unresolved MBHBs, will trace light MBHBs at very high redshift, placing constraints on black hole formation scenarios prior to the reionization epoch; the latter will be a formidable tool to follow the cosmic evolution of MBHs and the formation and dynamics of MBH binaries following galaxy mergers."916 The plan is as follows., The plan is as follows.917 In 2 we review the basics of the detection of CAV from MBIIDs. defining observable quantities such as the characteristic stram amplitude. signal-to-noise ratio. aud source detection rate.," In 2 we review the basics of the detection of GW from MBHBs, defining observable quantities such as the characteristic strain amplitude, signal-to-noise ratio, and source detection rate."918 In 3 we briefliv sunmnuarize our scenario for the cosmological evolution of ealaxy halos aud associated holes., In 3 we briefly summarize our scenario for the cosmological evolution of galaxy halos and associated holes.919 In £ we preseut confusion noise levels aud source uuuber counts., In 4 we present confusion noise levels and source number counts.920 Finally. in 85 we discuss our results.," Finally, in 5 we discuss our results."921" Following Thorne (1996). an interferometer can be characterized by two different sensitivitv curves. cepending ou the type of signal one expects to detect. 10. a ""burst? or a “periodic” GW source."," Following Thorne (1996), an interferometer can be characterized by two different sensitivity curves, depending on the type of signal one expects to detect, i.e. a “burst"" or a “periodic"" GW source."922" A burst. a short-lived VAignal whose waveform can be utterly complicated. cau be escribed in terms of a characteristic strain amplitude /, at the observed frequency f.c L/At.. where Af, is the uration of the signal (Thorne 1987)."," A burst, a short-lived signal whose waveform can be utterly complicated, can be described in terms of a characteristic strain amplitude $h_c$ at the observed frequency $f_c \sim 1/\Delta923t_s$ , where $\Delta t_s$ is the duration of the signal (Thorne 1987)."924 The spread of the power spectruni around £F. will be Af~f., The spread of the power spectrum around $f_c$ will be $\Delta f \sim f_c$.925 At the other xtreme. a perfectly periodic source cuits. for the cutire uration of the observation. at a fixed frequency f.," At the other extreme, a perfectly periodic source emits, for the entire duration of the observation, at a fixed frequency $f$."926 The power spectrum will be peaked at £. with a spread Af FN. where N is the umber of wave cveles clipped iuto he observation.," The power spectrum will be peaked at $f$, with a spread $\Delta f \simeq f/N$ , where $N$ is the number of wave cycles clipped into the observation."927 Iu this respect. a burst cau be thought as a sinele complete waveform with f£=f...," In this respect, a burst can be thought as a single complete waveform with $f=f_c$."928 In the case of a seriodic signal. the iuterferomieter sensitivity is increased w the fact that. across the observing interval r. the signal is repeated f7 times.," In the case of a periodic signal, the interferometer sensitivity is increased by the fact that, across the observing interval $\tau$, the signal is repeated $f\tau$ times."929 The seusitivitv to bursts (5) aud to periodic signals (hp) axe related by: Iu Fieure 1 the two curves fip aud δρ are colmpared for an assed Devear observation., The sensitivity to bursts $h_B$ ) and to periodic signals $h_P$ ) are related by: In Figure \ref{sensitivity} the two curves $h_B$ and $h_P$ are compared for an assumed 3-year observation.930 The curves are obtained combining the sinele-ain Michelson seusitivity curve (taken from the URL www.srlealtecliiedu/-shaue/seusitivitv) with the recent analysis of the instrumental noise below 10! Tz (Beuder 2003. extended from 3.«109 IIz to 1«10.5 Tz with a constant slope).," The curves are obtained combining the single-arm Michelson sensitivity curve (taken from the URL $\sim$ shane/sensitivity) with the recent analysis of the instrumental noise below $10^{-4}$ Hz (Bender 2003, extended from $3\times 10^{-6}$ Hz to $1\times 10^{-6}$ Hz with a constant slope)."931 Consider now a periodic signal of finite duration. with strain amplitude 5.," Consider now a periodic signal of finite duration, with strain amplitude $h$."932 The total energy carried by the wave will be proportional to the nuuiber of wave cycles à speut at that particular frequency., The total energy carried by the wave will be proportional to the number of wave cycles $n$ spent at that particular frequency.933" The quantity to be compared with 5j is then the ""characteristic strain Jh.—hn."," The quantity to be compared with $h_B$ is then the “characteristic"" strain $h_c \equiv h\sqrt{n}$ ."934 Note that for a periodic signal at frequency. f lasting for a time interval longer than the observation time 7. we have simply (9=fr.," Note that for a periodic signal at frequency $f$ lasting for a time interval longer than the observation time $\tau$ , we have simply $n=f\tau$."935 Then. the signal-to-noise ratio S/N increases by the same factor one would obtain comparing h to hp in equation (1)).," Then, the signal-to-noise ratio $S/N$ increases by the same factor one would obtain comparing $h$ to $h_P$ in equation \ref{eqburstperiodic}) )."936 The former approach. io. comparing ὃς top rather than / to fp. is more general. as it allows us to characterize the S/N not only for perfectly xeriodie signals (7= fr). or for bursts (0= 1). but also for events in which the emitted frequeucy shifts to iucreasiuglv arecr values during the spiral-in phase of the binary system.," The former approach, i.e. comparing $h_c$ to $h_B$ rather than $h$ to $h_P$, is more general, as it allows us to characterize the $S/N$ not only for perfectly periodic signals $n=f\tau$ ), or for bursts $n=1$ ), but also for events in which the emitted frequency shifts to increasingly larger values during the spiral-in phase of the binary system."937" Iu the latter case. 0=of) represeuts the nuuber of eveles spent in a frequency interval Af~f around yequency f£. aud hence fi, is the strain in a logarithnunic requeucyv interval (Flanagan Hughes 1998)."," In the latter case, $n=n(f)$ represents the number of cycles spent in a frequency interval $\Delta f \simeq f$ around frequency $f$, and hence $h_c$ is the strain in a logarithmic frequency interval (Flanagan Hughes 1998)."938 Typically. he timescale for frequency shift is long compared to he wave period. aud short compared to the duration of the observation.," Typically, the timescale for frequency shift is long compared to the wave period, and short compared to the duration of the observation."939 Onlv close to the innuenuoststable circular. orbit (CISCO). the CW frequency changes at a rate comparable to the frequency itself η~1 aud lence h.c h)," Only close to the innermoststable circular orbit (ISCO), the GW frequency changes at a rate comparable to the frequency itself $n\sim 1$ and hence $h_c \sim h$ )."940 Tu Figure dlowe also show 7L audbh. for two representative binary svstenis., In Figure \ref{sensitivity} we also show $h$ and$h_c$ for two representative binary systems.941" One should uote that the true observable CAV signal is. for f2ne/r (the ""kuec"," One should note that the true observable GW signal is, for $f> n/\tau$ (the “knee"""942of which stellar populations are detected at cach wavelength is essential to obtain the full picture.,of which stellar populations are detected at each wavelength is essential to obtain the full picture.943 Nevertheless. we are somewhat limited by the sensitivity. quality and field-of-view of the existing observations.," Nevertheless, we are somewhat limited by the sensitivity, quality and field-of-view of the existing observations."944 Star formation tends to be localised and varies within galaxies., Star formation tends to be localised and varies within galaxies.945 While the nuclear region ancl inner spiral arms of a galaxy are generally locations of significant star ormation. we also find new stars forming in other areas such as interaction zones and occasionally in isolated clumps (presumably. of high molecular gas density) in the zw outskirts of galaxies.," While the nuclear region and inner spiral arms of a galaxy are generally locations of significant star formation, we also find new stars forming in other areas such as interaction zones and occasionally in isolated clumps (presumably of high molecular gas density) in the far outskirts of galaxies."946 Vhe NGC 1512/1510 system is an xcellent. laboratory to study the locations ancl properties of its many star forming regions. from the galaxy. nuclei out o the largest racii where detached eclouds are found. (see Section 3.3) as well as in the interaction zone between the two galaxies.," The NGC 1512/1510 system is an excellent laboratory to study the locations and properties of its many star forming regions, from the galaxy nuclei out to the largest radii where detached clouds are found (see Section 3.3) as well as in the interaction zone between the two galaxies."947 llere we use a range of tracers to study the global SER of both NGC 1512 and. NGC 1510. (results. are summarised in Table 5. and Fig., Here we use a range of tracers to study the global SFR of both NGC 1512 and NGC 1510 (results are summarised in Table \ref{tab:sfr} and Fig.948 14). before investigating he local star formation activity within various parts of the GC 1512/1510 system (see Section 4.4).," 14), before investigating the local star formation activity within various parts of the NGC 1512/1510 system (see Section 4.4)."949 From our 20-cm radio continuum data we derive a recent global SER o£ Tor NGC 1512 and for NGC 1510. (see Section. 3.4)., From our 20-cm radio continuum data we derive a recent global SFR of for NGC 1512 and for NGC 1510 (see Section 3.4).950 Another. extinction-ree SER. estimate is derived. [rom the far-infrared. (£41) uminositv., Another extinction-free SFR estimate is derived from the far-infrared $FIR$ ) luminosity.951 Using the HtAS flux densities (Moshir οἱ al., Using the IRAS flux densities (Moshir et al.952 1990) together with the relations given by Sanders Alirabel (1996) and Ixennicutt (1998). we derive z012M.s [for NGC 1512 and [for NGC 1510.," 1990) together with the relations given by Sanders Mirabel (1996) and Kennicutt (1998), we derive $\approx$ for NGC 1512 and for NGC 1510."953 FIR emission comes from the thermal continuum re-radiation of dust. grains which absorb the visible and CY. radiation emitted by massive voung stars., $FIR$ emission comes from the thermal continuum re-radiation of dust grains which absorb the visible and $UV$ radiation emitted by massive young stars.954 In contrast. racio continuum emission is mainlv due to svnchrotron radiation from relativistic electrons. accelerated. in the remnants of core-collapse supernovae. therefore also associated with the presence of massive stars.," In contrast, radio continuum emission is mainly due to synchrotron radiation from relativistic electrons accelerated in the remnants of core-collapse supernovae, therefore also associated with the presence of massive stars."955 Both estimates trace the star formation activity in the last 100 Alvr., Both estimates trace the star formation activity in the last $\sim$ 100 Myr.956 However. as relativistic electrons have lifetimes of ~L00 Alvr (Conclon οἱ al.," However, as relativistic electrons have lifetimes of $\sim$ 100 Myr (Condon et al."957 2002). we should expect that the 20-em radio continuum emission traces SLRs with somewhat extended ages.," 2002), we should expect that the 20-cm radio continuum emission traces SFRs with somewhat extended ages."958 comission traces the most massive. ionising stars. ancl imescales of ~1O Myr. ic. the most recent events of star ormation in the galaxy.," emission traces the most massive, ionising stars, and timescales of $\sim$ 10 Myr, i.e. the most recent events of star formation in the galaxy."959 Lhe flux given by Meurer. ct al. (, The flux given by Meurer et al. (9602006) was corrected [for Galactic extinction. but not for internal extinction or for he contribution of the eemission lines adjacent to ((sec Lóppez-Sánnchez Esteban2008).,2006) was corrected for Galactic extinction but not for internal extinction or for the contribution of the emission lines adjacent to (see Lóppez-Sánnchez Esteban.961. Using the relation by Ixennicutt (1998). we find == 0.19 and for NGC 1512 and NGC 1510. respectively.," Using the relation by Kennicutt (1998), we find = 0.19 and for NGC 1512 and NGC 1510, respectively."962 Slighter lower values. == 0.13 and i. result when using the more recent Calzetti οἱ al. (," Slighter lower values, = 0.13 and , result when using the more recent Calzetti et al. ("9632007) calibration.,2007) calibration.964 UVemission. probes star formation over timescales of ~100 Mr. the life-time of the massive OB stars.," $UV$ -emission probes star formation over timescales of $\sim$ 100 Myr, the life-time of the massive OB stars."965 Using the extinetion-corrected CLALIZN. CV-magnitude. πριν. as eiven by Gil de Paz et al. (," Using the extinction-corrected GALEX $UV$ -magnitude, $m_{\rm FUV}$ , as given by Gil de Paz et al. ("9662007a). we derive the CV-MIux as follows: fpes teem 7 d = 140e10Dosohsft ,"2007a), we derive the $UV$ -flux as follows: $f_{\rm FUV}$ $^{-1}$ $^{-2}$ $^{-1}$ ] = $1.40 \times 10^{-15} \times 10^{0.4 \times (18.82 - m_{\rm FUV})}$."967We have Corrected mypvy for extinction assuming the Galactic value provided by Schlegel et al. (, We have corrected $m_{\rm FUV}$ for extinction assuming the Galactic value provided by Schlegel et al. (9681998). οV) = 0.011. and elpey=7.9EXG(D1.,"1998), $E(B-V)$ = 0.011, and $A_{\rm FUV} = 9697.9~E(B-V)$."970 Applying the Salim et al. (, Applying the Salim et al. (9712007) relation between the £UCV luminosity and the SER. we obtain == 0.12 and for NGC 1512 and NGC 1510. respectively.,"2007) relation between the $FUV$ luminosity and the SFR, we obtain = 0.12 and for NGC 1512 and NGC 1510, respectively."972 For comparison. applving the Ixennicutt (1998) relation results in values that are a 1.3 times higher.," For comparison, applying the Kennicutt (1998) relation results in values that are a 1.3 times higher."973 Here we prefer to use Salim οἱ al. (, Here we prefer to use Salim et al. (9742007) relation because it was derived. using CCALIZX data.,2007) relation because it was derived using GALEX data.975 The SINGS Legacy project lxennicutt et al., The SINGS Legacy project (Kennicutt et al.976 2003) provides Spitzer mic-infrarecd (ALL) images of GO 15191510., 2003) provides Spitzer mid-infrared $MIR$ ) images of NGC 1512/1510.977 ΑΙ emission. which traces the dust distribution. within galaxies. also agrees well with the »xosition of the C V-rieh star clusters in the svstem.," $MIR$ emission, which traces the dust distribution within galaxies, also agrees well with the position of the $UV$ -rich star clusters in the system."978 Because of its higher intrinsic. brightness. the MZ? emission. is mainlv detected. in the cores of both galaxies and in the inner ring of NGC 1512.," Because of its higher intrinsic brightness, the $MIR$ emission is mainly detected in the cores of both galaxies and in the inner ring of NGC 1512."979 Using the Spitzer 24jun. Hux density measurements of NGC 1512 (Dale et al., Using the Spitzer $\mu$ m flux density measurements of NGC 1512 (Dale et al.980 2007) and GC 1510 (obtained by us: see Table 5) together with the relations by Calzetti et al. (, 2007) and NGC 1510 (obtained by us; see Table 5) together with the relations by Calzetti et al. (9812007) we derive. SPisa = for NGC 1512 and forNGC 1510.,2007) we derive $SFR_{24\mu m}$ = for NGC 1512 and forNGC 1510.982" Combining the 24/2. luminosity (which.traces. the clust-absorbeck star formation) with the luminosity (which probes the unobscured. star formation) we derive. S£-Hg,", Combining the $\mu$ m luminosity (whichtraces the dust-absorbed star formation) with the luminosity (which probes the unobscured star formation) we derive $SFR_{\rm H\alpha+24\mu m}$ =983A system response function for both regions was needed to account for the contribution from the system setup (e.g. windows. beamsplitter. lens. filler aud detector) which alters the strength of the lines observed.,"A system response function for both regions was needed to account for the contribution from the system setup (e.g. windows, beamsplitter, lens, filter and detector) which alters the strength of the lines observed."984 This was achieved by recording the blackbody spectrum emitted from a solid graphite rod placed at the center of the Al3O tube: the end facing the speclrometer was machined into a concave cone to allow only blackbody emissions to be detected., This was achieved by recording the blackbody spectrum emitted from a solid graphite rod placed at the center of the $_{2}$ $_{3}$ tube; the end facing the spectrometer was machined into a concave cone to allow only blackbody emissions to be detected.985 The spectrum was then compared with a theoretical blackbody spectrum at the same temperature and normalized., The spectrum was then compared with a theoretical blackbody spectrum at the same temperature and normalized.986 The detector response curve for Region 2 and the effect on (he line intensities can be seenin Figure 4.., The detector response curve for Region 2 and the effect on the line intensities can be seenin Figure \ref{fig4}. .987"As an example. in Figure (4)) we consider the expected constraints on the BAO scale [rom a hypothetical galaxy survey with 10"" galaxies and bias of 1.5. spread over à volume of 0.8 Cpe? at small redshifts ?.","As an example, in Figure \ref{BAO_chi2}) ) we consider the expected constraints on the BAO scale from a hypothetical galaxy survey with $10^6$ galaxies and bias of 1.5, spread over a volume of 0.8 $^3$ at small redshifts ."988. In Figure (faa). we show the expected power spectrun--errors. normalized to a smooth power spectrum without BAO oscillations. using the fit to the transfer [unction.," In Figure \ref{BAO_chi2}a a), we show the expected power spectrum+errors, normalized to a smooth power spectrum without BAO oscillations, using the \citet{Eisenstein:1997jh} fit to the transfer function."989 The solid curve in Figure (4bb) shows NA? for fitting this spectrum with a different BAO (or sound horizon) scale. marginalizing over the amplitude of the oscillations.," The solid curve in Figure \ref{BAO_chi2}b b) shows $\Delta\chi^2$ for fitting this spectrum with a different BAO (or sound horizon) scale, marginalizing over the amplitude of the oscillations."990 Here. we only include linear scales. conservatively defined as &<0.1Mpe.t. and assume gaussian errors for ACh)—[|àPQ)/(222)]ol?.," Here, we only include linear scales, conservatively defined as $k<0.1 ~{\rm Mpc}^{-1}$, and assume gaussian errors for $\Delta(k) = \left[k^3P(k)/(2\pi^2)\right]^{1/2}$."991 While (he eaussian approximation to (he likelihood (dashed curve in Figure 4bb). is a good approximation-. close to the minimum.H itH grows indefinitely.B. while- the real A\>> saturates al c(S/N7.," While the gaussian approximation to the likelihood (dashed curve in Figure \ref{BAO_chi2}b b), is a good approximation close to the minimum, it grows indefinitely, while the real $\Delta\chi^2$ saturates at $\sim (S/N)^2$."992 In order to reflect this. we propose a simple analytic [function to approximate the (rue difference.B between 4L7 and ils. minimum. value: As shown in Figure (4bb). this takes the quadratie shape of the gaussian approximation close to the minimum since the denominator is then negligible. but Eq. ())," In order to reflect this, we propose a simple analytic function to approximate the true difference between $\chi^2$ and its minimum value: ^2 As shown in Figure \ref{BAO_chi2}b b), this takes the quadratic shape of the gaussian approximation close to the minimum since the denominator is then negligible, but Eq. \ref{dchi}) )"993 guarantees (hat A\? remains smaller than 65/.N)?. of the detection. [far from its minimum. which limits the statistical power of low signal-to-noise detections in constraining parameters.," guarantees that $\Delta\chi^2$ remains smaller than $(S/N)^2$ of the detection, far from its minimum, which limits the statistical power of low signal-to-noise detections in constraining parameters."994 While the interpolating function ()) is in good agreement with the actual A\7. we should note that it is only an approximation. and ideally one should use the full likelihood of the model fitting the data the [ull galaxy power spectrum) for an accurate statistical analvsis.," While the interpolating function \ref{dchi}) ) is in good agreement with the actual $\Delta\chi^2$, we should note that it is only an approximation, and ideally one should use the full likelihood of the model fitting the data the full galaxy power spectrum) for an accurate statistical analysis."995 Finally. Figure (5)) demonstrates (he effect of non-gaussian. posteriors on cosmological constraints.," Finally, Figure \ref{BAO_percival}) ) demonstrates the effect of non-gaussian posteriors on cosmological constraints."996 Here. we compare the eaussian approximation to likelihood «distribution [or distances (o 2=0.2 and :=0.35 in Sloan Digital Skv Survey 2010).. with our expectation from Equation ()).," Here, we compare the gaussian approximation to likelihood distribution for distances to $z=0.2$ and $z=0.35$ in Sloan Digital Sky Survey \citep[SDSS;][]{dr7}, with our expectation from Equation \ref{dchi}) )."997 Given that total S/N for BAO detection in Percivaletal.(2010) is ν10.1 for two degrees of [reedom.we see significant deviations from gaussian likelihoods bevond 99%. confidence level ?..," Given that total $S/N$ for BAO detection in \citet{dr7} is $\sqrt{13.1}$ for two degrees of freedom,we see significant deviations from gaussian likelihoods beyond $99\%$ confidence level ."998"As part of this effort to understand the nature of the evolution of dwarf galaxies, we present (Pilbratt 2010) photometric observation of the nearby (5.1 + 0.6 Mpc (Tosi et al.","As part of this effort to understand the nature of the evolution of dwarf galaxies, we present (Pilbratt 2010) photometric observation of the nearby (5.1 $\pm$ 0.6 Mpc (Tosi et al."999" 2001)) dwarf starburst galaxy, NGC 1705."," 2001)) dwarf starburst galaxy, NGC 1705."1000 The galaxy is dominated optically by a massive central super star cluster (SSC) NGC 1705-1 (Meurer et al., The galaxy is dominated optically by a massive central super star cluster (SSC) NGC 1705-1 (Meurer et al.1001" 1995), whilst studies in the mid and far-IR (Cannon et al."," 1995), whilst studies in the mid and far-IR (Cannon et al."1002 2006; Galametz et al., 2006; Galametz et al.1003" 2009) reveal the presence of two bright infrared regions flanking the central SSC, offset by ~250 pc from the SSC, with these off-nuclear regions dominating the global IR emission of NGC 1705."," 2009) reveal the presence of two bright infrared regions flanking the central SSC, offset by $\sim$ 250 pc from the SSC, with these off-nuclear regions dominating the global IR emission of NGC 1705."1004" This galaxy provides an ideal environment for exploring the effects of ongoing, massive star formation on the environment within a dwarf galaxy (Cannon et al."," This galaxy provides an ideal environment for exploring the effects of ongoing, massive star formation on the environment within a dwarf galaxy (Cannon et al."1005" 2006), given its sub-solar nebular metallicity (Z ~35% Zo; (Lee et al."," 2006), given its sub-solar nebular metallicity (Z $\sim$ $_{\odot}$; (Lee et al."1006 2004) and large reservoir of gas (Meurer et al., 2004) and large reservoir of gas (Meurer et al.1007" 1998), as we can trace the effects of the SSC on the surrounding interstellar medium, and in particular, can characterize the nature of dust and PAH emission."," 1998), as we can trace the effects of the SSC on the surrounding interstellar medium, and in particular, can characterize the nature of dust and PAH emission."1008 NGC 1705 was observed as part of the Dwarf Galaxy Survey programme (PI., NGC 1705 was observed as part of the Dwarf Galaxy Survey programme (PI.1009" S. Madden), a Guaranteed Time (GT) key program with the objective of mapping the dust and gas in 51 nearby dwarf galaxies, sampling a broad metallicity range of 1/50 to 1/3 Zo."," S. Madden), a Guaranteed Time (GT) key program with the objective of mapping the dust and gas in 51 nearby dwarf galaxies, sampling a broad metallicity range of 1/50 to 1/3 $_{\odot}$."1010 The galaxy was observed by SPIRE (Griffin et al., The galaxy was observed by SPIRE (Griffin et al.1011" 2010) at 250, 350 and 500 um for a total of 733 seconds."," 2010) at 250, 350 and 500 $\mu$ m for a total of 733 seconds."1012" in scan-map mode with scanning rate 30""/sec, with the final map covering roughly 16 x 16 arcmin."," in scan-map mode with scanning rate 30""/sec, with the final map covering roughly 16 x 16 arcmin."1013" The measured 1 o noise level are 5, 6 and 7 mJy beam""! at 250, 350 and 500 um respectively; the noise levels in the images are dominated by confusion."," The measured 1 $\sigma$ noise level are 5, 6 and 7 mJy $^{-1}$ at 250, 350 and 500 $\mu$ m respectively; the noise levels in the images are dominated by confusion."1014 The data were processed using the HIPE pipeline (see Pohlen et al. (, The data were processed using the HIPE pipeline (see Pohlen et al. (1015"2010) for a detailed description, Swinyard et al. (","2010) for a detailed description, Swinyard et al. ("10162010) for calibration accuracy and Bendo et al. (,2010) for calibration accuracy and Bendo et al. (10172010b) for details on the destriper).,2010b) for details on the destriper).1018" The pipeline produces maps with a pixel size of 6.0, 10.0 and 14.0"" at 250, 350 and 500 um respectively."," The pipeline produces maps with a pixel size of 6.0, 10.0 and 14.0"" at 250, 350 and 500 $\mu$ m respectively."1019 The ICC has released some interim small correction factors to improve the preliminary calibration., The ICC has released some interim small correction factors to improve the preliminary calibration.1020" All flux values derived using the current standard calibration file for the flux conversion, are multiplied by 1.02, 1.05, and 0.94, for the 250um, 350um,"," All flux values derived using the current standard calibration file for the flux conversion, are multiplied by 1.02, 1.05, and 0.94, for the $\mu$ m, $\mu$ m,"1021plane is 20°.,plane is $20\degr$.1022 The equatorial scattering region is assumed to be at least partly ionised and we therefore assume pure electron scattering with a Thomson optical depth of unity between the inner and outer radius of the wedge., The equatorial scattering region is assumed to be at least partly ionised and we therefore assume pure electron scattering with a Thomson optical depth of unity between the inner and outer radius of the wedge.1023 When showing modelling results we take into account existing (anti-)symmetries between the two hemispheres above and below the equatorial plane., When showing modelling results we take into account existing (anti-)symmetries between the two hemispheres above and below the equatorial plane.1024 We therefore only discuss the results for viewing directions 0?<i90° (or 1>cosi 0) with i being measured with respect to the symmetry axis of the disk and the torus., We therefore only discuss the results for viewing directions $0\degr < i < 90\degr$ (or $1 > \cos{i} > 0$ ) with $i$ being measured with respect to the symmetry axis of the disk and the torus.1025" The spectral modelling of the accretion disc irradiated by an elevated X-ray source confirms previous results, as is shown in Fig. 2.."," The spectral modelling of the accretion disc irradiated by an elevated X-ray source confirms previous results, as is shown in Fig. \ref{fig:disc}."1026" The reprocessed spectra reveal iron Ka and Kf fluorescence lines, the associated iron K absorption edge and the broad Comptonised hump centred around 30 keV. Absorption effects mainly due to hydrogen and helium produce an overall positive slope of the normalised reprocessed spectrum."," The reprocessed spectra reveal iron $\alpha$ and $\beta$ fluorescence lines, the associated iron K absorption edge and the broad Comptonised hump centred around 30 keV. Absorption effects mainly due to hydrogen and helium produce an overall positive slope of the normalised reprocessed spectrum."1027" Note that in the figure we have omitted the contribution of the directly visible, unpolarised primary radiation to have spectral features in the flux and polarisation spectrum come out more clearly."," Note that in the figure we have omitted the contribution of the directly visible, unpolarised primary radiation to have spectral features in the flux and polarisation spectrum come out more clearly."1028" In this modelling case, the polarisation vector, w, is always aligned with the projected symmetry axis."," In this modelling case, the polarisation vector, $\psi$, is always aligned with the projected symmetry axis."1029" At on view, when 7 is low, P does not exceed 1 per cent because the scattering geometry is almost symmetric."," At face-on view, when $i$ is low, $P$ does not exceed 1 per cent because the scattering geometry is almost symmetric."1030 The polarisation degree then rises with increasing 7 and the scattering medium appears less symmetric with respect to the line-of-sight., The polarisation degree then rises with increasing $i$ and the scattering medium appears less symmetric with respect to the line-of-sight.1031" The polarisation degree drops sharply across the iron K lines, which is due to dilution by the unpolarised fluorescent emission."," The polarisation degree drops sharply across the iron K lines, which is due to dilution by the unpolarised fluorescent emission."1032" Across the Compton hump the relation between P and i differs from the soft X-ray band, which is due to the Compton scattering phase function that favours forward over backward scattering at higher photon energies."," Across the Compton hump the relation between $P$ and $i$ differs from the soft X-ray band, which is due to the Compton scattering phase function that favours forward over backward scattering at higher photon energies."1033 It is instructive to compare these results to the non-relativistic calculations by Matt (1993).., It is instructive to compare these results to the non-relativistic calculations by \citet{matt1993}. .1034" Qualitatively, we find the same behaviour of the polarisation degree and angle but the absolute values of P obtained here are lower than for comparable cases shown in fig."," Qualitatively, we find the same behaviour of the polarisation degree and angle but the absolute values of $P$ obtained here are lower than for comparable cases shown in fig."1035 4 of Matt(1993)., 4 of \cite{matt1993}.1036. This can be explained when considering that the net polarisation degree of the reprocessed radiation results from integrating the reprocessed Stokes flux over the whole disc., This can be explained when considering that the net polarisation degree of the reprocessed radiation results from integrating the reprocessed Stokes flux over the whole disc.1037" After one scattering event, the polarisation of the outgoing radiation is directed perpendicularly to the scattering plane, and thus the central parts of the disc produce polarisation angles around i=0? (perpendicular to the symmetry axis) while the outer regions of the disc rather give rise to i)=90? (parallel to the projected symmetry axis)."," After one scattering event, the polarisation of the outgoing radiation is directed perpendicularly to the scattering plane, and thus the central parts of the disc produce polarisation angles around $\psi = 0\degr$ (perpendicular to the symmetry axis) while the outer regions of the disc rather give rise to $\psi = 90\degr$ (parallel to the projected symmetry axis)."1038 The net Stokes flux is dominated by the outer parts of the disc but still influenced also by the perpendicular polarisation state coming from the disc centre., The net Stokes flux is dominated by the outer parts of the disc but still influenced also by the perpendicular polarisation state coming from the disc centre.1039" Since Matt(1993) included a centralhole in the disc, the impact of the disc centre"," Since \cite{matt1993} included a centralhole in the disc, the impact of the disc centre"1040ido.,halo.1041 The latter would therelore give an even poorer fit than the spherical case. aud hence is 100 (θά here.," The latter would therefore give an even poorer fit than the spherical case, and hence is not tried here."1042 Therefore. in addition to the above parameters. (he axis ratio q is varied as well between 0.1 and 0.9 in steps of 0.1.," Therefore, in addition to the above parameters, the axis ratio $q$ is varied as well between 0.1 and 0.9 in steps of 0.1."1043 This gives a total of 47250 grid points to be scanned or each value of p., This gives a total of 47250 grid points to be scanned for each value of $p$.1044 We first thoroughly scan this grid to locate the region of mininnun V7., We first thoroughly scan this grid to locate the region of minimum ${\chi}^2$.1045 In retrospect. Naravan et al. (," In retrospect, Narayan et al. ("10462005) had pinned the rotation curve at a single point only ie the solar point with Ro — 8.5 kpc) using the local Oort constants “land D. for which the values are available for (he Galaxy.,"2005) had pinned the rotation curve at a single point only (i.e the solar point with R = 8.5 kpc) using the local Oort constants $A$ and $B$, for which the values are available for the Galaxy."1047 This effectively fixed the rotation curve locally. with respect to shape as well.," This effectively fixed the rotation curve locally, with respect to shape as well."1048 Also. the global trends exhibited by the observed curve was used as (he final criterion to choose the best fit density index (p = 2).," Also, the global trends exhibited by the observed curve was used as the final criterion to choose the best fit density index $p$ = 2)."1049 Here. on the other haud. we apply a more rigorous (treatinent by pinning the rotation curve at all the observed points.," Here, on the other hand, we apply a more rigorous treatment by pinning the rotation curve at all the observed points."1050 This. in fact. was imperative since the Oort A and £D constants for M31 are not known.," This, in fact, was imperative since the Oort $A$ and $B$ constants for M31 are not known."1051 For each of the above eric points. we evaluate (he galactic rotation curve using our (disc plus bulge plus halo) model as follows.," For each of the above grid points, we evaluate the galactic rotation curve using our (disc plus bulge plus halo) model as follows."1052" For an exponential disk. (he rotation velocity. 094,02) is given by (Dinnev Tremaine where My is the disk central surface density. £2; the disc scale leneth and yv. =R/2Ry. li being the galactocentric radius."," For an exponential disk, the rotation velocity $v_{disk}(R)$ is given by (Binney Tremaine where $\Sigma_{0}$ is the disk central surface density, $R_{d}$ the disc scale length and y =, R being the galactocentric radius."1053" J, and A, (where n=O and 1) are the modified. Bessel functions of the first aud second kind respectively.", $I_{n}$ and $K_{n}$ (where n=0 and 1) are the modified Bessel functions of the first and second kind respectively.1054 The above relation is for an infinitesimally (hin disk which we use here for simplicitv., The above relation is for an infinitesimally thin disk which we use here for simplicity.1055 For a thick disk. a separate result has to be used (as given in Beequaert Combes 1997). which we check gives à value within <14 of the value given by eq.(10). hence we are justified in using (he above simpler form.," For a thick disk, a separate result has to be used (as given in Becquaert Combes 1997), which we check gives a value within $< 1 \% $ of the value given by eq.(10), hence we are justified in using the above simpler form."1056 For the spherical bulge. rotation velocity Crag.(1?) is given by," For the spherical bulge, rotation velocity $v_{bulge}(R)$ is given by"1057"potential, we solve the nuclear network locally and then transfer generated energy on the multi-dimensional mesh.","potential, we solve the nuclear network locally and then transfer generated energy on the multi-dimensional mesh."1058 The transport of angular momentum may be also approximated by diffusion., The transport of angular momentum may be also approximated by diffusion.1059" Since the resultant distributions of thermodynamical quantities and elements will in general be non-uniform on the surface of constant effective potential, we will take their angular averages on the surface and solve the new rotational equilibrium for the obtained equations of state and rotation law."," Since the resultant distributions of thermodynamical quantities and elements will in general be non-uniform on the surface of constant effective potential, we will take their angular averages on the surface and solve the new rotational equilibrium for the obtained equations of state and rotation law."1060 This completes the single cycle and the iteration of this process will give the temporal evolution of rotational stars., This completes the single cycle and the iteration of this process will give the temporal evolution of rotational stars.1061 We hope that this procedure is feasible and that the formulation presented in this paper will contribute to the study of the influences of non-sphericity on the evolution of rapidly rotating massive stars., We hope that this procedure is feasible and that the formulation presented in this paper will contribute to the study of the influences of non-sphericity on the evolution of rapidly rotating massive stars.1062" Numerical computations were in part carried on XT4 and general common use computer system at the center for Computational Astrophysics, CfCA, the National Astronomical Observatory of Japan and on NEC-SX8 at Yukawa Institute for Theoretical Physics in Kyoto University."," Numerical computations were in part carried on XT4 and general common use computer system at the center for Computational Astrophysics, CfCA, the National Astronomical Observatory of Japan and on NEC-SX8 at Yukawa Institute for Theoretical Physics in Kyoto University."1063" This study was supported in part by the Grants-in-Aid for the Scientific Research from the Ministry of Education, Science and Culture of Japan (Nos."," This study was supported in part by the Grants-in-Aid for the Scientific Research from the Ministry of Education, Science and Culture of Japan (Nos."1064 80251403 and 19104006)., 80251403 and 19104006).1065iu Stokes I rom ?..,in Stokes I from \cite{hks+81}.1066" The simulator generates visibilities as they would be produced by the MWA correlator. :uxl the array parameters cal be couti«ος, by the user."," The simulator generates visibilities as they would be produced by the MWA correlator, and the array parameters can be controlled by the user."1067 In this case the caliration is perfect. tiere js Πο lonosplieric refraction. aud al the poi| sources and diffuse bacseround are intriisically uupolarized.," In this case the calibration is perfect, there is no ionospheric refraction, and all the point sources and diffuse background are intrinsically unpolarized."1068 The integration le weΡΕ λα)ely 2 hours aid ouly a siigle kkHz channel cenered at 160.02 MHz was sinulated., The integration time was approximately 2 hours and only a single kHz channel centered at 160.02 MHz was simulated.1069 The full obse‘vation has no been similated as the computatio1ial clemmauds are large aud a 'epresentaive inteeation is 'equi‘eck to test the fidelity of le Stokes conversion., The full observation has not been simulated as the computational demands are large and a representative integration is required to test the fidelity of the Stokes conversion.1070 The two hour itteeration is sullicie| 1ο enstre 1att je. doibait nolse SOLree in the simulated inap is from sidelobes., The two hour integration is sufficient to ensure that the dominant noise source in the simulated map is from sidelobes.1071 Tje linages were producecL by exactly tle salue pipeine as produced the real sky images. except that these 1ages we'e georeratect offine and not in 'eal time.," The images were produced by exactly the same pipeline as produced the real sky images, except that these images were generated offline and not in real time."1072 The FOV imaged was also slightly wicler: closer to ., The FOV imaged was also slightly wider: closer to $^\circ$.1073 AA has been peeled out he simulated images.e due to perfect calibratio there are no residtals 'emailning.," A has been peeled out the simulated images, due to perfect calibration there are no residuals remaining."1074 Figures 1. and 2 contaiu siguilicautly more poi sources than thes lmulation (Figure 7 )., Figures \ref{fig:StokesI_low} and \ref{fig:StokesI_mid} contain significantly more point sources than the simulation (Figure \ref{fig:mapsI}) ).1075 These sources are actually in the catalog., These sources are actually in the catalog.1076 However the siuation is very consealive when iichucling catalog sources. aud has only iucluded those soi‘ces With [lux meastwemeuts near the MWA observing baud aud that are predicted to be larger than ον.," However the simulation is very conservative when including catalog sources, and has only included those sources with flux measurements near the MWA observing band and that are predicted to be larger than Jy."1077 Nevertheless this sinulation gives a good indication of the performance of the systen under more controlled coucitious than the ineasured dataset., Nevertheless this simulation gives a good indication of the performance of the system under more controlled conditions than the measured dataset.1078 The RMS brightness of the polarization observatious is approximately of the Stoses { observatious. the point source polarized residual is cousicerably ess than thiS. of J0522-36027 at the very edge of the Q map. and only at the saue poln in the U uap.," The RMS brightness of the polarization observations is approximately of the Stokes I observations, the point source polarized residual is considerably less than this, of J0522–3627 at the very edge of the Q map, and only at the same point in the U map."1079 The residial is also a function of position «itlin the primary beam. near the center of the map the residual in QO is nearer and in the U 1laps there is almost no heasttlale association between Stoses { sources and polarized structure in the beam center - except as poarized residuals from tle peeling of AA. The polarized 'eslduals are a feature of üunperfect. calibration of the iicdiviclual antennas.," The residual is also a function of position within the primary beam, near the center of the map the residual in Q is nearer and in the U maps there is almost no measurable association between Stokes I sources and polarized structure in the beam center - except as polarized residuals from the peeling of A. The polarized residuals are a feature of imperfect calibration of the individual antennas."1080 The fact that the residuals deteriorate away from |Canu center lends further weight to thus claim: as differences betwee le primary beams become more apparent futher froin tlie beam center., The fact that the residuals deteriorate away from beam center lends further weight to this claim; as differences between the primary beams become more apparent further from the beam center.1081 The simulated point sotJyCes lave no intrinsic polarization aud as the calibration js perfect here should be no polarized j»oiut sources in tlie iutegrated maps. uulike the images of je true sky which do display some polarized point source sienals due to 1uperlect «alibration of le primary jeans.," The simulated point sources have no intrinsic polarization and as the calibration is perfect there should be no polarized point sources in the integrated maps, unlike the images of the true sky which do display some polarized point source signals due to imperfect calibration of the primary beams."1082 Figures 8 and 9 are featureless in this 'egard. with litle evideice of the exiseuce of any j»oint source features in the uaps.," Figures \ref{fig:mapsQ} and \ref{fig:mapsU} are featureless in this regard, with little evidence of the existence of any point source features in the maps."1083 The maps from simulation possess dMfuse polarization despite he fact that the iuput sky is uupolarized., The maps from simulation possess diffuse polarization despite the fact that the input sky is unpolarized.1084 This is due to tle polarized sidelobes o ‘the οokes ] emission., This is due to the polarized sidelobes of the Stokes I emission.1085 The polarized biehtuess in a elven pixel contaiis contribions from 1ot only the j»olarized sky at that position. but all the sideloyes of the οine and point source eiuissiou in the sky.," The polarized brightness in a given pixel contains contributions from not only the polarized sky at that position, but all the sidelobes of the diffuse and point source emission in the sky."1086 The sidelobes are generated with lustrumeual polarization couuneisurate with he position, The sidelobes are generated with instrumental polarization commensurate with the position1087"where #=1/3 if! <7, and &2-p/2 otherwise.",where $\kappa = 1/3$ if $\nu' < \nu'_m$ and $\kappa = -p/2$ otherwise.1088 Numerically speaking the integration procedure ts as follows., Numerically speaking the integration procedure is as follows.1089 First we tabulate A(t.) for à given set of physical parameters. so that we do not need to estimate it analytically but can use its exact dependence on the fluid Lorentz factor instead.," First we tabulate $R(t_e)$ for a given set of physical parameters, so that we do not need to estimate it analytically but can use its exact dependence on the fluid Lorentz factor instead."1090 We integrate over ( before we integrate over o. , We integrate over $\theta$ before we integrate over $\phi$ 1091is equal to the smoothing scale.,is equal to the smoothing scale.1092 Using cluster analysis we have detected two clusters in this sample a triplet and a quintet (8 out of a total of 21 quasars)., Using cluster analysis we have detected two clusters in this sample – a triplet and a quintet (8 out of a total of 21 quasars).1093 The z-sizes of both clusters are ~35h.1 Mpce., The $z$ -sizes of both clusters are $\sim 35h^{-1}$ Mpc.1094 The estimated. probability to be random is 0.05 and 0.01. respectively.," The estimated probability to be random is $0.05$ and $0.01$, respectively."1095 Both of these clusters contribute to the excess of quasar pairs at. separations of 30.40h! Ape., Both of these clusters contribute to the excess of quasar pairs at separations of $30-40h^{-1}$ Mpc.1096 ltesults for the BJS2 sample are shown in figs., Results for the BJS2 sample are shown in figs.1097 2c and d. The decay of BO) at ic<30h+ Mpe is caused by the Lact that the number density of quasars in this sample is considerably lower than in WI and (ZANDTB and even lower than in BIS] and BJS., 2c and d. The decay of $\Xi(r)$ at $r<30h^{-1}$ Mpc is caused by the fact that the number density of quasars in this sample is considerably lower than in KK and $^2$ B and even lower than in BJS1 and BJS3.1098 Εις sample therefore just lacks close pairs., This sample therefore just lacks close pairs.1099 The correlation Function Ductuates considerably at scales 1502005 Alpe (there is a significant excess of pairs at these separations as compared to the random distribution)., The correlation function fluctuates considerably at scales $\sim 150-200h^{-1}$ Mpc (there is a significant excess of pairs at these separations as compared to the random distribution).1100 These Huetuations correspond to the positive and negative peaks in AGU) at150h| Alpe and ~1805.7 Alpe., These fluctuations correspond to the positive and negative peaks in $\Delta\Theta(r)$ at $\sim 150h^{-1}$ Mpc and $\sim 180h^{-1}$ Mpc.1101 Using cluster analysis we have found a triplet (s-size ~205| Ape). à quartet (z-size 20h1 Alpe). and a sextet (z-size 75h Alpe).," Using cluster analysis we have found a triplet $z$ -size $\sim 20h^{-1}$ Mpc), a quartet $z$ -size $\sim 20h^{-1}$ Mpc), and a sextet $z$ -size $\sim 75h^{-1}$ Mpc)."1102 The probability that eachindividual cluster is random is smaller than 0.05., The probability that eachindividual cluster is random is smaller than $0.05$.1103 Our analysis has shown that the excess of pairs at separations 150200h Alpe is explained by the distance between the quartet and the sextet (Lsof1 Alpe)., Our analysis has shown that the excess of pairs at separations $150-200h^{-1}$ Mpc is explained by the distance between the quartet and the sextet $\sim 180h^{-1}$ Mpc).1104 The correlation function for this sample (lig., The correlation function for this sample (fig.1105 2e) has two broad “humps” at seales ~LOOM1 Alpe and ~2005.+ Alpe., 2e) has two broad “bumps” at scales $\sim 100h^{-1}$ Mpc and $\sim 200h^{-1}$ Mpc.1106 There are two negative peaks in the z;AO(r) corresponding to these fluctuations (Lig., There are two negative peaks in the $\Delta\Theta(r)$ corresponding to these fluctuations (fig.1107 21)., 2f).1108 Statistical significance of both peaks is 2a., Statistical significance of both peaks is $\sim 2\sigma$.1109 The cluster analysis failed to detect any large clusters in this sample., The cluster analysis failed to detect any large clusters in this sample.1110 However. the distribution of quasars in this sample is quite interesting.," However, the distribution of quasars in this sample is quite interesting."1111 Phere are seven relatively close pairs (distances between quasars in 5 of then are less than 305.+ Mpe and in the other two —407.+ Alpe) separated.olher by either SO.—1005.+ Alpe or 1802005+ Alpe., There are seven relatively close pairs (distances between quasars in 5 of them are less than $30h^{-1}$ Mpc and in the other two $\sim 40h^{-1}$ Mpc) separated by either $80-100h^{-1}$ Mpc or $180-200h^{-1}$ Mpc.1112 This causes the [luctuations in Z(r) at the corresponding scales., This causes the fluctuations in $\Xi(r)$ at the corresponding scales.1113 The results presented in the previous section suggest that the distribution of quasars in the analvzed. samples is not homogencous at scales of a [ον tens of megaparsecs., The results presented in the previous section suggest that the distribution of quasars in the analyzed samples is not homogeneous at scales of a few tens of megaparsecs.1114 Alany quasars belong to clumps of sizes 30TOh5 Alpe., Many quasars belong to clumps of sizes $30-70h^{-1}$ Mpc.1115 The clumps are often separated. hy 100200h4 Alpe. which creates a pair excess at the corresponding scales.," The clumps are often separated by $100-200h^{-1}$ Mpc, which creates a pair excess at the corresponding scales."1116 Qualitatively this quasar distribution is very similar to that of CIV. absorption svstenmis discussed. in the recent paper by Williger et al. (, Qualitatively this quasar distribution is very similar to that of CIV absorption systems discussed in the recent paper by Williger et al. (11171996).,1996).1118 They have found that their CIV sample contains two groups of 7 and 5 absorbers of sizes ~43511e69h7 Ape? and ~25453b% Alpe? (comoving) located at z2.3 and 2~2.5. respectively.," They have found that their CIV sample contains two groups of $7$ and $5$ absorbers of sizes $\sim43\times17\times69h^{-3}$ $^3$ and $\sim25\times4\times53h^{-3}$ $^3$ (comoving) located at $z\sim2.3$ and $z\sim2.5$, respectively."1119" The distance between these two groups (~501205+ Alpe) results in ""beating"" (the pair excess) giving rise to the correlation signal at these separations (3.5 significance level).", The distance between these two groups $\sim 50-120h^{-1}$ Mpc) results in “beating” (the pair excess) giving rise to the correlation signal at these separations $3.5\sigma$ significance level).1120 A number of smaller clumps were also detected., A number of smaller clumps were also detected.1121 The similar clusters of CIV. absorbers were also found in earlier studies bv Jakobsen Perryman (1992). Foltz et al. (," The similar clusters of CIV absorbers were also found in earlier studies by Jakobsen Perryman (1992), Foltz et al. ("11221993). and Dinshaw Impoey (1996).,"1993), and Dinshaw Impey (1996)."1123 Recently. Lespine Petitjean (1996) presented evidences for a coherent structure extende over ~SOP54 Alpe at zo2 in the distribution of meta absorption svstems.," Recently, Lespine Petitjean (1996) presented evidences for a coherent structure extended over $\sim80h^{-1}$ Mpc at $z\approx2$ in the distribution of metal absorption systems."1124 Although the numbers of CIV systems are also small. the similarity of the results may suggest tha both quasars ancl CIV. absorbers may. trace the same kin of underlying structures in the matter distribution.," Although the numbers of CIV systems are also small, the similarity of the results may suggest that both quasars and CIV absorbers may trace the same kind of underlying structures in the matter distribution."1125 Unlike Deng et al. (, Unlike Deng et al. (11261994) we did. not observe any evidence for a periodic signal in the function Or)..,1994) we did not observe any evidence for a periodic signal in the function $\Delta\Theta(r)$.1127. This may be cause bv the small number statistics., This may be caused by the small number statistics.1128 The clumps. distribution of quasars in the analvze samples is consistent. with the recent studies of the quasar distribution in the larger samples (Crampton. Cowley Hartwick LOST. 1989: Clowes Campusano 1991a. 1991b: Graham. Clowes Campusano 1995: Ixomboerg et al.," The clumpy distribution of quasars in the analyzed samples is consistent with the recent studies of the quasar distribution in the larger samples (Crampton, Cowley Hartwick 1987, 1989; Clowes Campusano 1991a, 1991b; Graham, Clowes Campusano 1995; Komberg et al."1129 1996)., 1996).1130 These samples were found to contain several relatively rich (~1025 QSOs) groups of quasars with sizes in the redshift clireetion of 70160b+ Mpe., These samples were found to contain several relatively rich $\sim10-25$ QSOs) groups of quasars with sizes in the redshift direction of $\sim70-160h^{-1}$ Mpc.1131 Ehe small extent of the pencil-bcam samples perpendicular to the line of sight prevents detection of such large groups., The small extent of the pencil-beam samples perpendicular to the line of sight prevents detection of such large groups.1132 However. the detected. smaller clumps can casily be parts of larger systems.," However, the detected smaller clumps can easily be parts of larger systems."1133 Lt would be very interesting to check this by studying larger deep samples which are currently. underway (e.g. Hall et al., It would be very interesting to check this by studying larger deep samples which are currently underway (e.g. Hall et al.1134 1996)., 1996).1135 If quasars anc CIV. absorption systems trace the matter distribution at high redshifts as galaxies or galaxy clusters co at. low redshifts. their clunmipy distribution suggests that. large-scale. inhomogeneitics similar to the nearby superclusters were already. distinct at 2—2.," If quasars and CIV absorption systems trace the matter distribution at high redshifts as galaxies or galaxy clusters do at low redshifts, their clumpy distribution suggests that large-scale inhomogeneities similar to the nearby superclusters were already distinct at $z\sim1-2$."1136 This information may provide some useful insights into the physics of high. redshift Universe., This information may provide some useful insights into the physics of high redshift Universe.1137 “Phe fact that we see structures at redshifts z~1.2 similar to the superclusters ab o2cO (dxomberg et al., The fact that we see structures at redshifts $z\sim1-2$ similar to the superclusters at $z\sim0$ (Komberg et al.1138 1996). for instance. favors low-clensity ACDAL or low-density CDM. models in. which perturbation amplitude at large. scales stops growing at 2cl.," 1996), for instance, favors low-density $\Lambda$ CDM or low-density CDM models in which perturbation amplitude at large scales stops growing at $z\geq 1$."1139 On the other hand. rather high ctuasar-quasar correlations at small separations and high number density contrasts in the detected quasar groups may indicate that o. distribution of quasars is highly biased. with respect to re matter distribution.," On the other hand, rather high quasar-quasar correlations at small separations and high number density contrasts in the detected quasar groups may indicate that the distribution of quasars is highly biased with respect to the matter distribution."1140 Although present available surveys we too small to provide a statistically reliable estimate X the power spectrum. 264). in the future. with bigger uasar samples and better models for both QSOs and CIV thsorbers. we will be able to ect useful constraints on the μα»eetrum. and thus on the theories of structure formation.," Although present available surveys are too small to provide a statistically reliable estimate of the power spectrum $P(k)$ , in the future, with bigger quasar samples and better models for both QSOs and CIV absorbers, we will be able to get useful constraints on the spectrum, and thus on the theories of structure formation,"1141varving between 36.000 ancl 000001 Gn steps of 100019) and AL varving over the [ull range ol the disk moclel grid. (here was no improvement in the fitting and the scale lactor-cderivecl distance was still grossly too large.,"varying between 36,000 and 60,000K (in steps of 1000K) and $\dot{M}$ varying over the full range of the disk model grid, there was no improvement in the fitting and the scale factor-derived distance was still grossly too large."1142 Next. we tested. (wo-lemperature fits with a cooler. slowly rotating photosphere and a hot. rapidly spinning. equatorial on as spe from disk accretion.," Next, we tested two-temperature fits with a cooler, slowly rotating photosphere and a hot, rapidly spinning, equatorial belt as expected from disk accretion."1143" In (his experiment. we varied the WD Τι"" between 42.000Ix. ancl 60.00Ix."" the accretion belt temperature between 50.000Ix. ancl 60.000Ix. in steps of 1000Ix. and the C and Si abundances of the white dwarl fixed."," In this experiment, we varied the WD $T_{eff}$ between 42,000K and 60,00K, the accretion belt temperature between 50,000K and 60,000K in steps of 1000K and kept the C and Si abundances of the white dwarf fixed."1144 Once again. as in the white dwarf plus disk case. there was no improvement in the 47 value.," Once again, as in the white dwarf plus disk case, there was no improvement in the $\chi^2_{\nu}$ value."1145 The quality of the model fits to the FUSE spectra of the (wo svstenis is «quite different., The quality of the model fits to the FUSE spectra of the two systems is quite different.1146 The fit to SS Aur is verv much in agreement with a model white dwarf atmosphere with log g — 9.0 and Tipp = 33.000Ix. This fit to the FUSE spectrum provides independent confirmation of the results of Lake&Sion(2001) who also found that the [ar UV IUE spectra were dominated bv a hot. massive white dwaif.," The fit to SS Aur is very much in agreement with a model white dwarf atmosphere with log $g$ = 9.0 and $T_{eff}$ = 33,000K. This fit to the FUSE spectrum provides independent confirmation of the results of \citet{lak01} who also found that the far UV IUE spectra were dominated by a hot, massive white dwarf."1147 The 7;;; they derived with IUE for the white dwarf in SS Aur was 30.000IX. .--This is surprising because it was widelv felt that the white dwarf in SS Aur was not the svstem was clisk-dominated in the lar UV. auc could not be analvzed unambiguously.," The $T_{eff}$ they derived with IUE for the white dwarf in SS Aur was 30,000K. This is surprising because it was widely felt that the white dwarf in SS Aur was not exposed, the system was disk-dominated in the far UV and could not be analyzed unambiguously."1148 In SS Aur. it dis also highly significant that there is little evidence of an additional hot component other (han a single temperature white dwar! photosphere.," In SS Aur, it is also highly significant that there is little evidence of an additional hot component other than a single temperature white dwarf photosphere."1149 The absence of (1175 À)) absorption in the FUSE spectrum suggests the possibility that the white dwarl is deficient in carbon., The absence of (1175 ) absorption in the FUSE spectrum suggests the possibility that the white dwarf is deficient in carbon.1150 If so. (his could be an indication (hat past thermonuclear processing (ancient novae) depleted the carbon.," If so, this could be an indication that past thermonuclear processing (ancient novae) depleted the carbon."1151 This possibility is supported by the indication that the N-abundance in the SS Aur WD surface lavers is elevated above solar., This possibility is supported by the indication that the N-abundance in the SS Aur WD surface layers is elevated above solar.1152 An alternative picture discussed by Gausickeetal.(2003). suggest that theN/C anomaly seen in the cbwarf novae BZ UMa. EY Cvg. IRNS J232953.9--062814. and now CIIUMa may have its origin ina CV with an originally more massive donor star (Alo>1.5M. ) which survived thermal time scale mass transfer (Schenkeretal.(2002) and references (herein).," An alternative picture discussed by \citet{gan03} suggest that the N/C anomaly seen in the dwarf novae BZ UMa, EY Cyg, 1RXS J232953.9+062814, and now CH UMa \citep{dul02,dul04} may have its origin in a CV with an originally more massive donor star $M_{2} > 1.5 M_{\odot}$ ) which survived thermal time scale mass transfer \citet{sch02} and references therein)."1153 In such asvstem. the white cdwarl would be accreting from the peeled away CNO-processed core stripped of its outer lavers during the (thermal timescale mass transfer.," In such a system, the white dwarf would be accreting from the peeled away CNO-processed core stripped of its outer layers during the thermal timescale mass transfer."1154 Ou FUSE spectrum of RU Peg likewise reveals à very hot white dwarf in agreement with the analvsis of the IUE archival spectra of RU Peg in quiescence., Our FUSE spectrum of RU Peg likewise reveals a very hot white dwarf in agreement with the analysis of the IUE archival spectra of RU Peg in quiescence.1155 We find that Τε = 49.000Ix for the white dwarf. is very close to the νε derived by Sion&Urban (2002)..," We find that $T_{eff}$ = 49,000K for the white dwarf, is very close to the $T_{eff}$ derived by \citet{sio02}. ."11562000 by FGMO03.,2000 by FGM03.1157 Outside of the flare the spectral parameters do not vary significantly and their average values (see below) are similar to the 2000 values (FGMO3). as is the X-ray luminosity (1.510 η).," Outside of the flare the spectral parameters do not vary significantly and their average values (see below) are similar to the 2000 values (FGM03), as is the X-ray luminosity $1.5\times 10^{30}$ )."1158 To derive the flaring emission spectral parameters. we applied the same procedure as for V827 Tau.," To derive the flaring emission spectral parameters, we applied the same procedure as for V827 Tau."1159 The quiescent spectral parameters are N(H)=2.43x107 επι. KT= keV. EM=140xI0? em. and Z=0.6Zs. and the resulting spectral parameters for the flaring component are listed in Table 5..," The quiescent spectral parameters are $=2.43 \times 10^{22}~{\rm cm^{-2}}$ , $kT=3.10$ keV, $E\!M= 1.401160\times 10^{53}$ $^{-3}$, and $Z=0.6~Z_{\odot}$, and the resulting spectral parameters for the flaring component are listed in Table \ref{tab:hl_flare}."1161 The event shows a peculiar evolution. with a monotonically decaying light curve associated with a highly irregular temperature evolution.," The event shows a peculiar evolution, with a monotonically decaying light curve associated with a highly irregular temperature evolution."1162 The temperature has two well-defined peaks above 7 keV separated by a deep minimum at about 3 keV. The temperature evolution suggests that this 1s the combination of two flares. probably physically related to each other but occurring in independent coronal structures.," The temperature has two well-defined peaks above 7 keV separated by a deep minimum at about 3 keV. The temperature evolution suggests that this is the combination of two flares, probably physically related to each other but occurring in independent coronal structures."1163 This kind of evolution has been predicted by modeling two independent flares by Reale.Güdel.Peres.&Audard(2004)., This kind of evolution has been predicted by modeling two independent flares by \cite*{rgp+2004}.1164. In order to support this hypothesis. we modeled the event by combining two flares computed with detailed hydrodynamie modeling of plasma confined in a coronal loop.," In order to support this hypothesis, we modeled the event by combining two flares computed with detailed hydrodynamic modeling of plasma confined in a coronal loop."1165 The light curve decay time suggests long flaring structures citealpsrj+91)). so that we considered each model flare to be identical to the one used to describe one of the flares observed during the COUP campaign (Favataetal.. 2005)) in detail.," The light curve decay time suggests long flaring structures \\citealp{srj+91}) ), so that we considered each model flare to be identical to the one used to describe one of the flares observed during the COUP campaign \citealp{ffr+2005}) ) in detail."1166 We assumed that each flare occurs in a coronal loop with à constant cross-section anc half-length L=107 em. symmetric around the loop apex.," We assumed that each flare occurs in a coronal loop with a constant cross-section and half-length $L = 10^{12}$ cm, symmetric around the loop apex."1167 Both flares were triggered by injecting a heat pulse in the loop. which was initially at à temperature of =20 MK.," Both flares were triggered by injecting a heat pulse in the loop, which was initially at a temperature of $\simeq 20$ MK."1168" This heat pulse is symmetrically deposited at the loop footpoints with a Gaussian spatial distribution of intensity 10 ergem ""land width 10!"" em (1/100 of the loop half-length).", This heat pulse is symmetrically deposited at the loop footpoints with a Gaussian spatial distribution of intensity 10 erg $^{-3}$ $^{-1}$ and width $10^{10}$ cm (1/100 of the loop half-length).1169 After 20 ks the heat pulse was switched off completely., After 20 ks the heat pulse was switched off completely.1170 From the evolution of the plasma density and temperature along the loop computed with the Palermo-Harvard hydrodynamic loop model (Peresetal..1982.. Bettaetal..1997)). we synthesized the corresponding EPIC spectra of the loop throughout the flare. deriving a light curve and the evolution of temperature.," From the evolution of the plasma density and temperature along the loop computed with the Palermo-Harvard hydrodynamic loop model \citealp{psv+82}, \citealp{bpr+97}) ), we synthesized the corresponding EPIC spectra of the loop throughout the flare, deriving a light curve and the evolution of temperature."1171 To model the HL Tau event we duplicated. the resulting light curve and temperature evolution with à time shift.," To model the HL Tau event we duplicated, the resulting light curve and temperature evolution with a time shift."1172 The two flares are identical flares. except for a normalization factor. which represents the loop cross-section and does not enter explicitly in the hydrodynamic modeling.," The two flares are identical flares, except for a normalization factor, which represents the loop cross-section and does not enter explicitly in the hydrodynamic modeling."1173 The second flare has a normalization factor of 0.3 aa correspondingly smaller cross-section) and it starts 60 ks after the first., The second flare has a normalization factor of 0.3 a correspondingly smaller cross-section) and it starts 60 ks after the first.1174 We summed the resulting two asynchronous sequences of flare spectra and obtained a single sequence of spectra. which we integrated to derive a single light curve and fit with single temperature EPIC model spectra.," We summed the resulting two asynchronous sequences of flare spectra and obtained a single sequence of spectra, which we integrated to derive a single light curve and fit with single temperature EPIC model spectra."1175 Figure 14 shows the resulting light curve and temperature evolution as compared to those obtained from the data., Figure \ref{fig:hlflare} shows the resulting light curve and temperature evolution as compared to those obtained from the data.1176 The model temperatures in the first flare are somewhat higher than the observed one. but the main flare characteristics. tthe monotonic light curve and the temperature dip. are reproduced by the double-flare model well — although the model is not unique.," The model temperatures in the first flare are somewhat higher than the observed one, but the main flare characteristics, the monotonic light curve and the temperature dip, are reproduced by the double-flare model well – although the model is not unique."1177 Since no constraint can be derived from the data. each of the two flares was modeled with no significant residual heating present during the flare decay citealprbp+97)).," Since no constraint can be derived from the data, each of the two flares was modeled with no significant residual heating present during the flare decay \\citealp{rbp+97}) )."1178 Such modeling implies very large flaring structures. similar to the ones found in ONC YSOs by Favataetal.(2005) — where the data allowed investigation of the presence of sustained heating.," Such modeling implies very large flaring structures, similar to the ones found in ONC YSOs by \cite{ffr+2005} – where the data allowed investigation of the presence of sustained heating."1179 Such large structures. with L= 5R.. have only been found in YSOs. and were interpreted by Favataetal.(2005) as linking the star to the aceretion disk. aas being the magnetic structures supporting the magnetospheric accretion.," Such large structures, with $L \simeq11805\,R_*$ , have only been found in YSOs, and were interpreted by \cite{ffr+2005} as linking the star to the accretion disk, as being the magnetic structures supporting the magnetospheric accretion."1181 In addition to the evidence from the ONC YSOs. HL Tau is the first Taurus YSO in which such large flaring structures have been detected.," In addition to the evidence from the ONC YSOs, HL Tau is the first Taurus YSO in which such large flaring structures have been detected."1182 We cannot a priori exclude that shorter loops with sustained heating in the decay may also reproduce the features of this flare., We cannot a priori exclude that shorter loops with sustained heating in the decay may also reproduce the features of this flare.1183 However. the long delay of the model flares required to reproduce the distant temperature peaks suggests that very large structures must be involved in the flare (Realeetal.. 2004)).," However, the long delay of the model flares required to reproduce the distant temperature peaks suggests that very large structures must be involved in the flare \citealp{rgp+2004}) )."1184 Unlike the other stars in the present study. 2285845 (an active binary system) is not à member of the star-forming region. on the basis of its radial velocity and. proper motion (Walteretal.. 1988)).," Unlike the other stars in the present study, 285845 (an active binary system) is not a member of the star-forming region, on the basis of its radial velocity and proper motion \citealp{wbm+88}) )."1185 The primary spectral type is G8. and Schneideretal.(1998) report a separation of 73 mas and a magnitudedifference of 1.19 mag.," The primary spectral type is G8, and \cite{shw98}1186 report a separation of 73 mas and a magnitudedifference of 1.19 mag."1187 In the 2000 observation. 2285845 showed significant variability. and similar behavior is present in. the present data set (Fig. 15))," In the 2000 observation, 285845 showed significant variability, and similar behavior is present in the present data set (Fig. \ref{fig:hdall}) )."1188 Table 12— summarizes the best-fit, Table \ref{tab:hd_all11pn} summarizes the best-fit1189"The joint distribution in z and. /. in presenceof the M—L luminosity (equation (16)]) ls The mean red-shift is The number clensitv of galaxies per unit [Iux interval is The number of galaxies as given bv formula (38)) has à maximum al 2,4, where A comparison between the observed and theoretical number of galaxies as given by the M—-L [function is reported in Figure 13. where the 2dF Galaxy. Redshift Survey is considered ancl in Figure 14 where the 6dE Galaxy Survey is considered.","The joint distribution in $z$ and $f$, in presenceof the ${\mathcal M}-L$ luminosity (equation \ref{equation_schechter_mia}) )) is The mean red-shift is The number density of galaxies per unit flux interval is The number of galaxies as given by formula \ref{nfunctionz_mia}) ) has a maximum at $z_{max}$ where A comparison between the observed and theoretical number of galaxies as given by the ${\mathcal M}-L$ function is reported in Figure \ref{maximum} where the 2dF Galaxy Redshift Survey is considered and in Figure \ref{maximum_6d} where the 6dF Galaxy Survey is considered."1190 One method to deduce the mass of a star by its absolute visual magnitude is presented: ihe mass of a galaxy is deduced by analogy., One method to deduce the mass of a star by its absolute visual magnitude is presented; the mass of a galaxy is deduced by analogy.1191 In the case of the galaxies. the bolometric correction of the stus will be replaced bv the sun's absolute magnitude and mass-Inninosity ratio different in each selected band.," In the case of the galaxies, the bolometric correction of the stars will be replaced by the sun's absolute magnitude and mass-luminosity ratio different in each selected band."1192 In the case of the stars il is possible to parameterise (he mass of the star. My . as a [unction of the observable colour (23—V). see Zaninetli (2005)..," In the case of the stars it is possible to parameterise the mass of the star , ${\mathcal M_S}$ , as a function of the observable colour $(B-V)$, see \cite{zaninetti05} ."1193. The first equation connectsthe (5—V) colour with the temperature, The first equation connectsthe $(B-V)$ colour with the temperature1194»220 ms7!. of which 43 exhibit significant periodicities.,"20 m$^{-1}$, of which 43 exhibit significant periodicities."1195coniplicate processes is far bevoud the scope of this Letter.,complicate processes is far beyond the scope of this Letter.1196 Nonetheless. the actual flux may uot be far dinuuer than our prediction since averagelv speaking. the LFs of the slow u-shells aud the i-shells are significantly different. but their masses are comparable.," Nonetheless, the actual flux may not be far dimmer than our prediction since averagely speaking, the LFs of the slow n-shells and the i-shells are significantly different, but their masses are comparable."1197 So the interaction of them can power a bright UV flash iu the range ~Lot1050., So the interaction of them can power a bright UV flash in the range $\sim 10^{13}-10^{15}{\rm cm}$.1198 Y. Z. F thanks T. Lu aud Z. Li for their long-term chcouragcinent on the subject of ueutrou-fed CRBs., Y. Z. F thanks T. Lu and Z. Li for their long-term encouragement on the subject of neutron-fed GRBs.1199 We also thauk the anomvmous referee for herμας coustructive conunents., We also thank the anonymous referee for her/his constructive comments.1200 This work is supported by the National Natural Scicuce Foundation (erauts 1052325011 and 10233010). the National 973 Project ou Fundamental Researches of China (NINBRSF C19990751).," This work is supported by the National Natural Science Foundation (grants 10225314 and 10233010), the National 973 Project on Fundamental Researches of China (NKBRSF G19990754)."

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