ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 We can then define a (local) adiabatic exponent through the identification P?=αρ., We can then define a (local) adiabatic exponent through the identification $P=\kappa \rho^{\gamma}$.3 If none of the equilibrium quantities depend on the height z: or angle 6. we can take Fourier transforms in the axial ancl azimuthlial directions.," If none of the equilibrium quantities depend on the height $z$ or angle $\theta$, we can take Fourier transforms in the axial and azimuthal directions."4" Neelecting (he perturbed gravitational potential 0,. the equations for the normal modes of Lagrangian perturbations. (£=£(r)e/mÜwl) )"," Neglecting the perturbed gravitational potential $\delta \Phi_g$ , the equations for the normal modes of Lagrangian perturbations, $\xxi=\xxi(r)e^{ i(k_z z+m \theta -\omega t)} $ )"5 2)).,\ref{table_lensing}) ).6" With the full extended model, the effect is especially noticeable on N, and β, showing that there are significant degeneracies between the effect of neutrinos on the matter power spectrum and the effect a scale-dependent primordial power spectrum with several degrees of freedom."," With the full extended model, the effect is especially noticeable on $N_\nu$ and $\beta$, showing that there are significant degeneracies between the effect of neutrinos on the matter power spectrum and the effect a scale-dependent primordial power spectrum with several degrees of freedom."7 Fig., Fig.8" 2 (red bars) shows that the greatest degradation in constraints with respect to QCDM occurs in theparameters Wa, h, and ns."," \ref{hist_lensing} (red bars) shows that the greatest degradation in constraints with respect to QCDM occurs in theparameters $w_a$, $h$, and $n_s$."9 There is an additional degeneracy in the matter power spectrum between the small-scale power-suppression effect of massive neutrinos and the form of the primordial power spectrum for certain values of the primordial spectral index., There is an additional degeneracy in the matter power spectrum between the small-scale power-suppression effect of massive neutrinos and the form of the primordial power spectrum for certain values of the primordial spectral index.10" With weak lensing only, we obtain tighter constraints on α than on ns, which is in agreement with the results obtained by Kitchingetal.(2008a) and Ishaketal.(2004)."," With weak lensing only, we obtain tighter constraints on $\alpha$ than on $n_s$ , which is in agreement with the results obtained by \citet{Kitching2008} and \citet{Ishak:2004}."11". Although the precision for both parameters is degraded when neutrinos are added, this error hierarchy is preserved, even when CMB constraints are added (see section 3.5 below)."," Although the precision for both parameters is degraded when neutrinos are added, this error hierarchy is preserved, even when CMB constraints are added (see section \ref{joint} below)."12 The addition of priors has a significant effect on parameter constraints., The addition of priors has a significant effect on parameter constraints.13" The FoM is improved by a factor of 6 for the vQCDM+a+8 model (Table 3)), and we obtain better constraints for all parameters, especially Qb, Opg (related to the geometry of the universe), h and ns."," The FoM is improved by a factor of 6 for the $\nu\mr{QCDM}+\alpha+\beta$ model (Table \ref{Table_Joint}) ), and we obtain better constraints for all parameters, especially $\Omega_b$, $\Omega_\mr{DE}$ (related to the geometry of the universe), $h$ and $n_s$."14" Adding CMB priors also lifts the degeneracy between some parameters, so the extension of the parameter set does not significantly degrade the error bars."," Adding CMB priors also lifts the degeneracy between some parameters, so the extension of the parameter set does not significantly degrade the error bars."15" This can be seen in Fig. 4,,"," This can be seen in Fig. \ref{hist_joint},"16" where Q,, h and ns, which are well-constrained byPlanck, are now hardly affected by the addition of extra parameters."," where $\Omega_b$, $h$ and $n_s$, which are well-constrained by, are now hardly affected by the addition of extra parameters."17" It can be seen from Table 3 that we obtain better constraints on a than on ns with the addition of CMB priors, reversing the error hierarchy obtained with lensing only."," It can be seen from Table \ref{Table_Joint} that we obtain better constraints on $\alpha$ than on $n_s$ with the addition of CMB priors, reversing the error hierarchy obtained with lensing only."18" Moreover, the addition of neutrino parameters does not significantly affect the precision on ns."," Moreover, the addition of neutrino parameters does not significantly affect the precision on $n_s$."19" With combined calculations, we obtain an improvement in the joint (wo,wa) constraints (Fig. 5))."," With combined calculations, we obtain an improvement in the joint $w_0,\,w_a$ ) constraints (Fig. \ref{ellipses_article_joint}) )."20 The constraints are robust against the addition of neutrino parameters and the primordial power spectrum parameters o and f., The constraints are robust against the addition of neutrino parameters and the primordial power spectrum parameters $\alpha$ and $\beta$.21 In this section we study the effect of the survey design on our error forecasts., In this section we study the effect of the survey design on our error forecasts.22" The optimisation for an all-sky tomographic weak lensing survey has been investigated by (Amara&Réfrégier 2007),, who use the dark energy FoM as the optimisation benchmark."," The optimisation for an all-sky tomographic weak lensing survey has been investigated by \citep{AR2007}, , who use the dark energy FoM as the optimisation benchmark."23" Our survey configuration is determined by the parameters: the area Ας, the median redshift zm, and the observed number density of galaxies ng."," Our survey configuration is determined by the parameters: the area $A_s$, the median redshift $z_m$, and the observed number density of galaxies $n_g$."24 The lensing correlation function is additionally defined by the range of multipoles over which it is measured., The lensing correlation function is additionally defined by the range of multipoles over which it is measured.25" In order to investigate the dependence of the marginalized errors in our /QCDM-4-aB parameter set on the survey parameters, we calculate the lensing Fisher matrix while varying one survey parameter at a time."," In order to investigate the dependence of the marginalized errors in our $\nu\mr{QCDM}+\alpha+\beta$ parameter set on the survey parameters, we calculate the lensing Fisher matrix while varying one survey parameter at a time."26 In Fig., In Fig.27" 6 we show the relative marginalized error, defined as Ap/|p| for different values of the median redshift zm and the maximum multipole Cmax."," \ref{opt} we show the relative marginalized error, defined as $\Delta p/ |p|$ for different values of the median redshift $z_m$ and the maximum multipole $\ell_\mr{max}$."28" We note that the scaling with z,, is similar for all the parameters in our three sectors of interest, with N, and B showing a stronger dependence on z,;,."," We note that the scaling with $z_m$ is similar for all the parameters in our three sectors of interest, with $N_\nu$ and $\beta$ showing a stronger dependence on $z_m$."29" We find that all parameters have a roughly similar scaling with €max, and that the greatest gain in precision is observed in the range 10?<fmax«10*."," We find that all parameters have a roughly similar scaling with $\ell_\mr{max}$, and that the greatest gain in precision is observed in the range $10^2<\ell_\mr{max}<10^4$."30" We also carried out the same calculation for various values of the survey area A, and the galaxy count ng, finding a linear scaling between the parameter precision and these two survey parameters."," We also carried out the same calculation for various values of the survey area $A_s$ and the galaxy count $n_g$, finding a linear scaling between the parameter precision and these two survey parameters."31" Our results show that the survey area has the greatest effect on parameter precision, and are in agreement with Amara&Réfrégier(2007)."," Our results show that the survey area has the greatest effect on parameter precision, and are in agreement with \citet{AR2007} ."32" Our results show that the optimum survey strategy holds not just for the joint dark energy parameters (wo,wa), but also for the other parameters in the cosmological model."," Our results show that the optimum survey strategy holds not just for the joint dark energy parameters $(w_0,w_a)$, but also for the other parameters in the cosmological model."33" A different survey design would therefore lead to a rescaling of the marginalized errors shown in Table 2,, but would not significantly modify the results shown in Fig. 2.."," A different survey design would therefore lead to a rescaling of the marginalized errors shown in Table \ref{table_lensing}, but would not significantly modify the results shown in Fig. \ref{hist_lensing}."34 This indicates that the effect of extending the hypothesis space is independent of the survey design., This indicates that the effect of extending the hypothesis space is independent of the survey design.35" The main aim of this paper was to put three different sectors of the cosmological model: dark energy, dark matter (cold and neutrinos) and initial conditions on equal footing while forecasting constraints for a future weak lensing survey."," The main aim of this paper was to put three different sectors of the cosmological model: dark energy, dark matter (cold and neutrinos) and initial conditions on equal footing while forecasting constraints for a future weak lensing survey."36" To do this we introduce a new parameter 8, which models the second order running spectral index."," To do this we introduce a new parameter $\beta$, which models the second order running spectral index."37" We have forecast errors for an all-sky tomographic weak lensing survey, and for weak lensing--Planck, using different cosmological parameter sets, studying the effect of the addition of parameters in the model."," We have forecast errors for an all-sky tomographic weak lensing survey, and for weak , using different cosmological parameter sets, studying the effect of the addition of parameters in the model."38" We have shownthat error forecasts for some parameters are stable against changes in the parameter set (Table 2)), and that degeneracies between the dark energy parameters wo and uw, arenot significantly affected by"," We have shownthat error forecasts for some parameters are stable against changes in the parameter set (Table \ref{table_lensing}) ), and that degeneracies between the dark energy parameters $w_0$ and $w_a$ arenot significantly affected by"39gaalaxies.,alaxies.40A‘nett.Meakin.&You1gsr(2009) [οιthat QALL Was Lol coustéult. but. depen προι the flow properties. aid the equation state.,"\cite{amy09} foundthat $\aml$ was not constant, but depended upon the flow properties, and the equation of state."41 For solar models tle surface convectic zoue is deep. aud changes ittle. so taking a οομσδίαμί aaL is an aclecuate appronimaticdL for his particular exaripe.," For solar models the surface convection zone is deep, and changes little, so taking a constant $\alpha_{ML}$ is an adequate approximation for this particular example."42 Ju MILT. he velocity obtaine by a «'ouvective eddy is computed frcxu the work cloie by the ΡΟον force over a lini& lenetn where D is the comp'essibilitv.. Ap the pressure scale height. auk AVVv-Vad is the usual “super-acliaalle excess.”," In MLT, the velocity obtained by a convective eddy is computed from the work done by the buoyancy force over a mixing length: where $\beta_T$ is the compressibility, $H_P$ the pressure scale height, and $ \Delta \nabla \equiv \nabla - \nabla_{ad} $ is the usual “super-adiabatic excess.”"43 For a given convective luniiosity. larger aa implies larger velocities.," For a given convective luminosity, larger $\aml$ implies larger velocities."44 Shallow cOuvection ZOLes. having shorter distances fo* buoyant acc'eleratjon 1ο work. will havesmaller values of QALL and smaller velocity scaes (Arnett.leakiu.&Your18oO 2009).," Shallow convection zones, having shorter distances for buoyant acceleration to work, will havesmaller values of $\aml$ and smaller velocity scales \citep{amy09}."45. As te depth of tle convection zone increases. tL esize of the larges eddies. also rises. implyiOm€ arger Qayp-," As the depth of the convection zone increases, the size of the largest eddies also rises, implying larger $\aml$."46 Such au iucrease will not coninte indefinitely: more vigorous convection «eveOps ore violeu clissipatioli., Such an increase will not continue indefinitely; more vigorous convection develops more violent dissipation.47" The value of QajL seenis to ""saliübale"" for very deep coivection zones (Arielt.Meakili.&Young20()):leakinArnet 2009)."," The value of $\aml$ seems to “saturate” for very deep convection zones \citep{amy09,ma09}."48". The solar conve""ion zone is 20 pressure scale ieiglits. dee>» and has yet to be simulated or its full «lepth with resolution as high as sed in Me:du&Arnett(2007) or Stein&Norcdlui (1995)."," The solar convection zone is 20 pressure scale heights deep, and has yet to be simulated for its full depth with resolution as high as used in \cite{ma07b} or \cite{sn98}."49. Here we will examine e Case i which such saturatiou occurs OALLcAv| , Here we will examine the case in which such saturation occurs at $ \aml \approx 4$.50This may be app‘opriate [or e situulatious of Nordlund aix Stein (R. ein. privae communication) aud those of eakin&Arnett(2009).. and 1s consistent the inseusitivity of the Stein&Nord-(1998) simulations o the exact position of the lower boundary. which was deeper han this.," This may be appropriate for the simulations of Nordlund and Stein (R. Stein, private communication) and those of \cite{ma09}, and is consistent with the insensitivity of the \cite{sn98} simulations to the exact position of the lower boundary, which was deeper than this."51 Fu‘ther analysis of 1lis issue is progress \feakin&Arnett(2009):: 3D Mlations for COLVECive zoues of depth 0.5 2 pressure scales heights seem cousistent ith this inte‘pretation., Further analysis of this issue is in progress \cite{ma09}; 3D simulations for convective zones of depth 0.5 to 5 pressure scales heights seem consistent with this interpretation.52 The distribution of ues for ασε in Tab e2 covers this range., The distribution of values for $\aml$ in Table \ref{tablep} covers this range.53 Solter ectrations of state. such as in partial lonlzation ZOLLES or electrou-positron pair Zones. give less vigorous velocities. but do uot change the qualitative picture (A1ell.Meakin.&Young 2009).," Softer equations of state, such as in partial ionization zones or electron-positron pair zones, give less vigorous velocities, but do not change the qualitative picture \citep{amy09}."54.. The simul:idOLS ol Porter&Woodward (2000).. for an ideal gas equatlor1 of state. also seein {ὁ σιeoestao that saturation may be beginuing a‘ound QaprLP2 3. which is consistent.," The simulations of \cite{pw00}, , for an ideal gas equation of state, also seem to suggest that saturation may be beginning around $\aml \approx 3$ , which is consistent."55 Figure 1 shows the evolutionary tracks, Figure \ref{fighr} shows the evolutionary tracks56galaxy to reach the background red continuum on either side of the lline.,galaxy to reach the background red continuum on either side of the line.57 In Fig., In Fig.58" 1 we illustrate for all the galaxies, maps of the continuum-subtracted iintensity, vvelocity, and vvelocity dispersion (σ) maps over the entire field-of-view as indicated in Table 1.."," \ref{fig:allmaps} we illustrate for all the galaxies, maps of the continuum-subtracted intensity, velocity, and velocity dispersion $\sigma$ ) maps over the entire field-of-view as indicated in Table \ref{tab:sampleparams}."59" For 3342, the ddata only cover the central ~2’ rradius, and only the mmapsare displayed here."," For 342, the data only cover the central $\approx 2$ radius, and only the mapsare displayed here."60 Corresponding B-band images from the Digitized Sky Survey (DSS) are also presented for comparison with the Fabry-Perot maps., Corresponding $B$ -band images from the Digitized Sky Survey (DSS) are also presented for comparison with the Fabry-Perot maps.61" We quantify the observed velocity fields by using the tilted-ring method combined with the harmonic decomposition formalism (e.g., ??).."," We quantify the observed velocity fields by using the tilted-ring method combined with the harmonic decomposition formalism \citep[e.g., ][]{Schoenmakersetal1997, Fathietal2005}. ."62 We assume that, We assume that63eas (Table 3)).,gas (Table \ref{table_xrayfit3}) ).64 The spectrum may be fit bv abundance sets with and without iron. the other elements being adjusted in a plausible fashion. and approximately matching the abundances found for the nebular gas and the stellar atmosphere.," The spectrum may be fit by abundance sets with and without iron, the other elements being adjusted in a plausible fashion, and approximately matching the abundances found for the nebular gas and the stellar atmosphere."65 If an abundance set depleted in iron is adopted. the {flux predicted may be brought into agreement with: observations. but (his also serves to illustrate that il is inescapable (hat iron is depleted in the plasma emitüng the N-ravs inBD+30°3639.," If an abundance set depleted in iron is adopted, the flux predicted may be brought into agreement with observations, but this also serves to illustrate that it is inescapable that iron is depleted in the plasma emitting the X-rays in."66. Clearly. it would be exiremely useful to determine both the atmospheric iron abundance of the central star as well as (he nebular iron abundance. since these would set bounds on the abundances expected in the hot gas.," Clearly, it would be extremely useful to determine both the atmospheric iron abundance of the central star as well as the nebular iron abundance, since these would set bounds on the abundances expected in the hot gas."67 The depletion of iron in the gas emilling in A-ravs is a strong chemical signature that is of potential use in identibving the origin of this gas., The depletion of iron in the gas emitting in X-rays is a strong chemical signature that is of potential use in identifying the origin of this gas.68 The eas emitting in N-ravs should have the composition of either the nebular gas. the stellar wind. or some mixture of the two.," The gas emitting in X-rays should have the composition of either the nebular gas, the stellar wind, or some mixture of the two."69 Nebular iron abundances are rare., Nebular iron abundances are rare.70 IIowever. (μον do exist for NGC 6543 and NGC 7027.," However, they do exist for NGC 6543 and NGC 7027."71 Perinottoetal(1999). found depletions by [actors of 11 and 80 in NGC 6543 and NGC 1027. respectively. (he extreme values thev found lor their sample of four planetary nebulae.," \citet{perinottoetal1999} found depletions by factors of 11 and 80 in NGC 6543 and NGC 7027, respectively, the extreme values they found for their sample of four planetary nebulae."72 Sinlarly. Sterlingοἱal(2005) found a depletion of iron by a factor of 3-14 in SwSt-l depending upon the region of the nebula observed.," Similarly, \citet{sterlingetal2005} found a depletion of iron by a factor of 3-14 in SwSt-1 depending upon the region of the nebula observed."73 It is therefore clear that iron maa often be depleted by a [actor of 10 or more in planetiry nebulae. though: more extensive studies would obviously be verv helpful.," It is therefore clear that iron may often be depleted by a factor of 10 or more in planetary nebulae, though more extensive studies would obviously be very helpful."74 The remaining fraction of the iron is presumably found in dust grains., The remaining fraction of the iron is presumably found in dust grains.75 Similar depletion of iron onto dust grains is also found in 1I regions. e.g... Rodriguez(2002) found iron depletions by [actors of 3 to 50.," Similar depletion of iron onto dust grains is also found in H regions, e.g., \citet{rodriguez2002} found iron depletions by factors of 3 to 50."76 Unfortunately. nothing is known of the iron abundance in the stellar wind for the objects studied here.," Unfortunately, nothing is known of the iron abundance in the stellar wind for the objects studied here."77 aand NGC 6543 have hyvdrogen-celicient central stars. NGC 7009 has a hvdrogen-rich central star. and nothing is known of the central star in NGC 7027.," and NGC 6543 have hydrogen-deficient central stars, NGC 7009 has a hydrogen-rich central star, and nothing is known of the central star in NGC 7027."78 Generally. there appears to be a clear trend (hat hydrogen-deficent central stars as well as their supposed PG-1159 aud WO progeny. with or without winds. have a(mospheres deficient in iron by (vpically at least an order of magnitude (Miksaetal.2002:Herald&Bianchi2004a.b.c:Stasiiskaal2004).," Generally, there appears to be a clear trend that hydrogen-deficient central stars as well as their supposed PG-1159 and WO progeny, with or without winds, have atmospheres deficient in iron by typically at least an order of magnitude \citep{miksaetal2002, heraldbianchi2004a, heraldbianchi2004b, heraldbianchi2004c,stasinskaetal2004}."79. This iron deficiency may be the result of nuclear processing (e.g..Herwigetal.2003).," This iron deficiency may be the result of nuclear processing \citep[e.g.,][]{herwigetal2003}."80. On the other hand. the central stus of planetary nebulae with hyvdrogen-vich atmospheres that still have winds do not appear to be depleted in iron (Herald&Bianchi2004a.b.c)..," On the other hand, the central stars of planetary nebulae with hydrogen-rich atmospheres that still have winds do not appear to be depleted in iron \citep{heraldbianchi2004a, heraldbianchi2004b, heraldbianchi2004c}."81 The depletion of iron in either the stellar wind or nebular materialcould explain our, The depletion of iron in either the stellar wind or nebular materialcould explain our82Starburst (SB) galaxies are known to have had an important role in. galaxy formation and evolution throughout the whole history of the Universe.,Starburst (SB) galaxies are known to have had an important role in galaxy formation and evolution throughout the whole history of the Universe.83 In the far-infrared (FIR) and sub-millimeter. observations are limited by extragalactic confusion: details on small spatial scales are lost in the noise because of the poor angular resolution of the instruments.," In the far-infrared (FIR) and sub-millimeter, observations are limited by extragalactic confusion: details on small spatial scales are lost in the noise because of the poor angular resolution of the instruments."84 As a result. unresolved starburst galaxies form the cosmic infrared background (CIB) (2?2).. which peaks at around 200 jan. In the mid-infrared. a large fraction of the CIB has been resolved into individual sources: ?).. for instance. resolved 70 of the 24 jim background.," As a result, unresolved starburst galaxies form the cosmic infrared background (CIB) \citep{1996A&A...308L...5P,1998ApJ...508..123F, 1999A&A...344..322L}, which peaks at around 200 $\mu$ m. In the mid-infrared, a large fraction of the CIB has been resolved into individual sources: \citet{2004ApJS..154...70P}, for instance, resolved 70 of the 24 $\mu$ m background."85 In the FIR. before the advent of the Herschel telescope. a smaller fraction has been resolved: with Spitzer. ?) resolved and of the CIB at 70 anc 160 um. respectively.," In the FIR, before the advent of the Herschel telescope, a smaller fraction has been resolved: with Spitzer, \citet{2004ApJS..154...87D} resolved and of the CIB at 70 and 160 $\mu$ m, respectively."86 2). managed to resolve of the CIB at 70 jm) using a very deep but small field. which was thus limited by cosmic variance.," \citet{2006ApJ...647L...9F} managed to resolve of the CIB at 70 $\mu$ m using a very deep but small field, which was thus limited by cosmic variance."87 More recently. ?) integratec counts coming from Herschel/PACS data at 100 and 160 jim and resolved ~45% and ~52% of the CIB. respectively.," More recently, \citet{2010AA...518L..30B} integrated counts coming from Herschel/PACS data at 100 and 160 $\mu$ m and resolved $\sim45\%$ and $\sim52\%$ of the CIB, respectively."88 At longer wavelengths. ?) directly resolved15%..1O%.. ane of the CIB at 250. 350. and 500 um.respectively. using Herschel/SPIRE data.," At longer wavelengths, \citet{2010AA...518L..21O} directly resolved, and of the CIB at 250, 350, and 500 $\mu$ m,respectively, using Herschel/SPIRE data."89 Confusion can be circumvented by the use of statistical methods., Confusion can be circumvented by the use of statistical methods.90 For instance. by stacking 24 jm sources. ?) were able to resolve a large fraction of the CIB at 70 jim and 160 jm and ?) increased their fractions from 45% to 50%. and from 52% to 75% at 100 pm and 160 jm. respectively.," For instance, by stacking 24 $\mu$ m sources, \citet{2006A&A...451..417D} were able to resolve a large fraction of the CIB at 70 $\mu$ m and 160 $\mu$ m and \citet{2010AA...518L..30B} increased their fractions from $\%$ to $\%$, and from $\%$ to $\%$ at 100 $\mu$ m and 160 $\mu$ m, respectively."91 Using a P(D) approach. ?) were able to obtain still larger fractions of and89%.. at 100 anc 160 um. respectively; P(D) derived counts of ?) account for 64. 60. and of the CIB at 250. 350. and 500 jum. respectively.," Using a P(D) approach, \citet{2011AA...532A..49B} were able to obtain still larger fractions of and, at 100 and 160 $\mu$ m, respectively; P(D) derived counts of \citet{2010MNRAS.409..109G} account for 64, 60, and of the CIB at 250, 350, and 500 $\mu$ m, respectively."92 These results imply that the sources detectec at 24 yam constitute the bulk of the CIB around its peak., These results imply that the sources detected at 24 $\mu$ m constitute the bulk of the CIB around its peak.93 ?) showed that galaxies that dominate the emission at 24 jum become more and more luminous and massive as the redshift increases starting from luminous infrared galaxies (LIRGs) with 10!ZL.<Lig10Li. at 0.8<z1.2 with intermediate mass. to ultra-luminous infrared galaxies (ULIRGs) with Οἱ].<Lg10L. that dominate at ς>2 and have masses >10!Ms. The clustering of galaxies that make up the CIB can be characterized by its anisotropies.," \citet{2006A&A...454..143C} showed that galaxies that dominate the emission at 24 $\mu$ m become more and more luminous and massive as the redshift increases starting from luminous infrared galaxies (LIRGs) with $10^{11}L_{\odot}<L_{IR}<10^{12}L_{\odot}$ at $0.8<z<1.2$ with intermediate mass, to ultra-luminous infrared galaxies (ULIRGs) with $10^{12}L_{\odot}<L_{IR}<10^{14}L_{\odot}$ that dominate at $z>2$ and have masses $> 10^{11}M_{\odot}$ The clustering of galaxies that make up the CIB can be characterized by its anisotropies."94 This clustering was. first detected at 160 uim with Spitzer (??).. and then measured at 250. 350. and 500 yam using BLAST data (?)..," This clustering was first detected at 160 $\mu$ m with Spitzer \citep{2007ApJ...665L..89L, 2007A&A...474..731G}, and then measured at 250, 350, and 500 $\mu$ m using BLAST data \citep{2009ApJ...707.1766V}."95 All three of these sets of data enabled the detection of an excess of signal on intermediate spatial scales caused by the clustering of starburst galaxies which enabled them to derive the linear bias parameter 5 that relates the density fluctuations of Iuminous matter to those of dark matter (DM)., All three of these sets of data enabled the detection of an excess of signal on intermediate spatial scales caused by the clustering of starburst galaxies which enabled them to derive the linear bias parameter $b$ that relates the density fluctuations of luminous matter to those of dark matter (DM).96 ?) measured a value of b=2.44£0.2 while ?) obtained b=3+0.3., \citet{2007ApJ...665L..89L} measured a value of $b=2.4\pm0.2$ while \citet{2009ApJ...707.1766V} obtained $b=3\pm0.3$.97 The difference may be due to selection effects., The difference may be due to selection effects.98 At longer wavelengths. higher redshift SB galaxies are probed (?).. 2))) and at these higher redshifts. SB galaxies are a highly biased tracer of the underlying dark matter density field.," At longer wavelengths, higher redshift SB galaxies are probed \cite{2005ARA&A..43..727L}, \cite{2008A&A...481..885F}) ) and at these higher redshifts, SB galaxies are a highly biased tracer of the underlying dark matter density field."99 They indeed formed in very massive DM halos early in the history of the Universe., They indeed formed in very massive DM halos early in the history of the Universe.100 ?) derived the two-point correlation function of 24 pm selected sources divided into two redshift bins (0.6<z1.2 and zo» 1.6). finding that these SB galaxies are strongly clustered and embedded in DM halos of =10'M.. for the high z sample and =10!!5A£. for the low = one.," \citet{2008MNRAS.383.1131M} derived the two-point correlation function of 24 $\mu$ m selected sources divided into two redshift bins $0.6<z<1.2$ and $z>1.6$ ), finding that these SB galaxies are strongly clustered and embedded in DM halos of $\simeq10^{12.8}M_{\odot}$ for the high $z$ sample and $\simeq10^{11.8}M_{\odot}$ for the low $z$ one."101 ?) computed the angular correlation function with Herschel/SPIRE data., \citet{2010AA...518L..22C} computed the angular correlation function with Herschel/SPIRE data.102" They found that 250 um sources are embedded in DM halos of ~10At, at <z >~2.1. whereas bright 500 pm sources reside in more massive halos. ~10M. at <z >~2.6."," They found that 250 $\mu$ m sources are embedded in DM halos of $\sim10^{12}M_{\odot}$ at $<z>\sim$ 2.1, whereas bright 500 $\mu$ m sources reside in more massive halos, $\sim10^{13}M_{\odot}$, at $<z>\sim$ 2.6."103 The CIB anisotropy measurements in the FIR and submillimeter were followed by those of ?) at 1.3 mm and 2 mm with the South Pole Telescope and by the ?) at 1.4 and 2 mm with the Atacama Cosmology Telescope (?)))., The CIB anisotropy measurements in the FIR and submillimeter were followed by those of \citet{2010ApJ...718..632H} at 1.3 mm and 2 mm with the South Pole Telescope and by the \citet{2010AAS...21538407F} at 1.4 and 2 mm with the Atacama Cosmology Telescope \citet{2010AAS...21538408D}) ).104 More recently. ?) derived CIB power spectra from to simultaneously at 350. 550. 850. and 1380 jm in six high-Galactic latitude fields.," More recently, \citet{2011A&A...536A..18P} derived CIB power spectra from to simultaneously at 350, 550, 850, and 1380 $\mu$ m in six high-Galactic latitude fields."105 ?) extended the measurements to smaller angular scales. using Herschel/SPIRE observations of the Lockman-hole field at 250. 350. and 500 jm. These measurements allow us to start to refine the analysis of the clustering properties of galaxies responsible for the CIB. and its cosmic evolution to high redshift (z~ ," \citet{2011Natur.470..510A} extended the measurements to smaller angular scales, using Herschel/SPIRE observations of the Lockman-hole field at 250, 350, and 500 $\mu$ m. These measurements allow us to start to refine the analysis of the clustering properties of galaxies responsible for the CIB, and its cosmic evolution to high redshift $\sim$ "106The above criteria to determine the amount of gas in cach category (cooled or both cold aud dense) are applied ou a cell by cell basis within the virial radius of cach peak in the data sample.,The above criteria to determine the amount of gas in each category (cooled or both cold and dense) are applied on a cell by cell basis within the virial radius of each peak in the data sample.107 The temperature threshold has been chosen to eusure that the gas is substantially cooler than the virial temperature of the peak as given in Equation 5.., The temperature threshold has been chosen to ensure that the gas is substantially cooler than the virial temperature of the peak as given in Equation \ref{eq:mtz}.108 The density threshold for cold. dense eas corresponds to the gas density at which the barvous ecole miportaut to the eravitational poteutial aud thus o the subsequent evolution of the core.," The density threshold for cold, dense gas corresponds to the gas density at which the baryons become important to the gravitational potential and thus to the subsequent evolution of the core."109 However. setting a densitythreshold for eas that has been able to cool through he chemistry but may not he dense is more subtle.," However, setting a density threshold for gas that has been able to cool through the chemistry but may not be dense is more subtle."110 We uust he careful o exclude eas that is cool only because it belongs to lowey qmass clouds that are infalliug outo the nore luassive structure., We must be careful to exclude gas that is cool only because it belongs to lower mass clouds that are infalling onto the more massive structure.111 With the high mass resolution in our simulations. this substructure due to merging is resolvable.," With the high mass resolution in our simulations, this substructure due to merging is resolvable."112 Since there is uo cooling in the peaks extracted roni the simulation without.. aux cold gas preseut uust be due to iufall and mereiue in the outer regions of he cloud.," Since there is no cooling in the peaks extracted from the simulation without, any cold gas present must be due to infall and merging in the outer regions of the cloud."113 We coufirin that this is indeed true in the data set without cooling by looking at radial profiles for the cold gas fraction in several peaks., We confirm that this is indeed true in the data set without cooling by looking at radial profiles for the cold gas fraction in several peaks.114 We then used this data ο determine the density threshold for cooled gas given above that first muninuzes the cold iufalliug componcut., We then used this data to determine the density threshold for cooled gas given above that first minimizes the cold infalling component.115 As shown in the bottom panel of Figure 10.. the resultaut vackeround due to merging is uceligible for all but the uost massive peaks m the sample.," As shown in the bottom panel of Figure \ref{fig:coldfraction}, the resultant background due to merging is negligible for all but the most massive peaks in the sample."116" The bump in f, iu the Heh mass end of the bottom paucl is due to the merecr of the two most massive peaks iu the box at 2—20.5 seen in Fieure &..", The bump in $f_c$ in the high mass end of the bottom panel is due to the merger of the two most massive peaks in the box at $z \sim 20.5$ seen in Figure \ref{fig:mvir_z}.117 Even in this extreme case the above density hreshold keeps the infalliug cold component of f£. luted o a few percent., Even in this extreme case the above density threshold keeps the infalling cold component of $f_c$ limited to a few percent.118 Note also that £4. the fraction of gas hat can become cease aud form stars. ix zero for all peaks iu the simulation without.. as expected.," Note also that $f_{cd}$, the fraction of gas that can become dense and form stars, is zero for all peaks in the simulation without, as expected."119 Figure 10. represeuts one of the key results of this paper., Figure \ref{fig:coldfraction} represents one of the key results of this paper.120 Both the fraction of eas f£. that can cool aud the fraction fea of cold. dense gas available for star formation merease ogarithinically with the preealactic cloud mass over the nass range dominated by cooling.," Both the fraction of gas $f_c$ that can cool and the fraction $f_{cd}$ of cold, dense gas available for star formation increase logarithmically with the pregalactic cloud mass over the mass range dominated by cooling."121 This depeudenuce can be summarized in a simple fitting formula: or (|=ο (cooled gas} or /=ed (cold. deuse gas) and AMc Atty.," This dependence can be summarized in a simple fitting formula: for $i=c$ (cooled gas) or $i=cd$ (cold, dense gas) and $M > M_{TH}$ ."122" The best fit slopes D. for the fraction of eas f. that cools are 0.138+0.006. 0.139+0.008. and 1.085x0.009 for radiative background fluxes of Fri:=0. 10??, and 10?3+... respectively. where he quoted errors are standard errors of the fit."," The best fit slopes $B_c$ for the fraction of gas $f_c$ that cools are $0.138 123\pm 0.006$, $0.139 \pm 0.008$, and $0.085 \pm 0.009$ for radiative background fluxes of $F_{LW}= 0$, $10^{-22}$, and $10^{-21}$, respectively, where the quoted errors are standard errors of the fit."124 However. he fraction of cold. deuse gas f; iucreases much less rapidly with mass.," However, the fraction of cold, dense gas $f_{cd}$ increases much less rapidly with mass."125" Best fit slopes μμ for this cold. dense component are 0.058c0.006. 0.066+0.007. and 1.030+0.001 for radiative background fluxes Fy=0. 10??, aud 107)|. respectively."," Best fit slopes $B_{cd}$ for this cold, dense component are $0.058 \pm 0.006$, $0.066 \pm 0.007$, and $0.030 \pm 0.004$ for radiative background fluxes $F_{LW}= 0$, $10^{-22}$, and $10^{-21}$, respectively."126 The equality of the slopes for flux levels Jpy=0 aud 10.72 is striking and probably reflects the fact that the cooling chemistry is the same in each case independent of radiative flux level., The equality of the slopes for flux levels $J_{LW} = 0$ and $10^{-22}$ is striking and probably reflects the fact that the cooling chemistry is the same in each case independent of radiative flux level.127 The appareut softening of the slopes for Fri=1075 tis most likely an artifact of the small nuniber of independent high mass clouds in our smiulatiou volume., The apparent softening of the slopes for $F_{LW} = 10^{-21}$ is most likely an artifact of the small number of independent high mass clouds in our simulation volume.128 The mass thresholds AMrg for nonzero f. and fi ave about equal at fixed Fr: and increase as the level of soft UV flux is iucreased., The mass thresholds $M_{TH}$ for nonzero $f_c$ and $f_{cd}$ are about equal at fixed $F_{LW}$ and increase as the level of soft UV flux is increased.129 This isa characterization of the negative feedback expected due to the photodissociation of iu the lower mass clouds., This is a characterization of the negative feedback expected due to the photodissociation of in the lower mass clouds.130 Our results for the mass threshold can be parameterized as The preseuce of a soft photodissociating flux at the levels considered here delays. but does not prevent eas in these pregalactic objects from becoming deuse due to cooling in their cores.," Our results for the mass threshold can be parameterized as: The presence of a soft photodissociating flux at the levels considered here delays, but does not prevent gas in these pregalactic objects from becoming dense due to cooling in their cores."131 The fraction of gas available for star formation due to cooling can be represeuted as: with Mr given by Equation 8.., The fraction of gas available for star formation due to cooling can be represented as: with $M_{TH}$ given by Equation \ref{eq:mth}.132 Iu this section we use spherically averaged racial profiles of the pregalactie cloud properties to illustrate the internal dynamics of its collapse., In this section we use spherically averaged radial profiles of the pregalactic cloud properties to illustrate the internal dynamics of its collapse.133 We first consider the evolution with redslift of a collapsing preegalactic object when exposed to a fixed level of plotodissociative flux., We first consider the evolution with redshift of a collapsing pregalactic object when exposed to a fixed level of photodissociative flux.134 We then consider how the internal properties of a given cloud change at fixed redshift when the level of the photodissociating flux changes., We then consider how the internal properties of a given cloud change at fixed redshift when the level of the photodissociating flux changes.135 Iu three of the panels of Figure rofüe:POevolve we show the evolution of the eas density. teiiperature. and mass fraction from 2=30° to maximal refinement at 2=21.5 for the first object to collapse at a radiative-. flux level Fry=E10D7 i. (," In three of the panels of Figure \\ref{fig:P0evolve} we show the evolution of the gas density, temperature, and mass fraction from $z=30$ to maximal refinement at $z=21.5$ for the first object to collapse at a radiative flux level $F_{LW}=10^{-21}$ . ("136See Table 1)),See Table \ref{tab:first}) ).137 In the fourth panel (lower right) we show the evolution of the cooling aud dyaunical timescales for 25<2<21.5 after cooling beeius., In the fourth panel (lower right) we show the evolution of the cooling and dynamical timescales for $25 \le z \le 21.5$ after cooling begins.138" Starting at large radius the upper set of curves represent the cooling time £f,"" given by where sg. ay. aud sg aro nunibor deusities for the eax. neutral hwdrogen. aud respectively, Ayis the Boltzmann constant. Fis the temperature. and A is the cooling function."," Starting at large radius the upper set of curves represent the cooling time $t_{cool}$ given by where $n_g$, $n_H$, and $n_{H2}$ are number densities for the gas, neutral hydrogen, and respectively, $k_B$is the Boltzmann constant, $T$ is the temperature, and $\Lambda$ is the cooling function."139" The lower set of curves at large radii show the dvuamical tiuescale f,5, where forthe same redshifts.", The lower set of curves at large radii show the dynamical timescale $t_{dyn}$ where forthe same redshifts.140 The solid (dot-dashed) horizoutal line is the IIubble time, The solid (dot-dashed) horizontal line is the Hubble time141speed of Alfvóuu velocity. ο=(B72/1zp)7.,"speed of Alfvénn velocity, $v_A=(B^2/4\pi\rho)^{1/2}$."142 The global magnetic field dissipation time should be not shorter thu {σαχ, The global magnetic field dissipation time should be not shorter than $l/v_A$.143 Notice that couverting a siguificant power into X-ravs needs a Thomson optical depth rr of orders of 1., Notice that converting a significant power into X-rays needs a Thomson optical depth $\tau_T$ of orders of 1.144" Putting all these together. one vields the time scale of variability: where B,=B/10! Gauss. AL|;—AL/(OOP exes/s)."," Putting all these together, one yields the time scale of variability: where $B_4=B/10^4$ Gauss, $\Delta L_{45}=\Delta L/(10^{45}$ ergs/s)."145 Af is sensitive to maguetic field streugth. but only weakly depends on the luuinositv.," $\Delta t$ is sensitive to magnetic field strength, but only weakly depends on the luminosity."146 For the fastest variation cing the rug phase of the flare. AL=5«10!! cres/s iu Af=280 sec (using the linear fit result). one vields B>EN10477? Gauss for the flare reeion.," For the fastest variation during the rising phase of the flare, $\Delta L=5\times 10^{44}$ ergs/s in $\Delta t= 280$ sec (using the linear fit result), one yields $B\ge 5\times 10^4 \tau_T^{3/8}$ Gauss for the flare region."147 If part of the energy dissipated via magnetie reconnection is converted into kinetic euerev. causing bulk motion ofthe faring material (Beloborodoy 1999).," If part of the energy dissipated via magnetic reconnection is converted into kinetic energy, causing bulk motion of the flaring material (Beloborodov 1999)."148 The bulk motion would boost the apparent variability if it is towards the observer., The bulk motion would boost the apparent variability if it is towards the observer.149 However. the flare material is only mildly relativistic. as estimated by Beloborodov(1999). this would not seriously affect the magnetic field given above.," However, the flare material is only mildly relativistic, as estimated by Beloborodov(1999), this would not seriously affect the magnetic field given above."150 The strength of the magnetic field in an accretion disk is in principle Iuuited by equipartition with the disk pressure., The strength of the magnetic field in an accretion disk is in principle limited by equipartition with the disk pressure.151 Stronger magnetic field will rise buovautlv from the accretion disk. leading to a magueticallv confined corona (Galeev. Rosner Vaiana 1979).," Stronger magnetic field will rise buoyantly from the accretion disk, leading to a magnetically confined corona (Galeev, Rosner Vaiana 1979)."152 Mineshige et al. (, Mineshige et al. (1532000) argued that the maeuetic field is large in the NLSIs due to aree pressure caused by the trapped pliotous iu the iuner region (radiation pressure dominated) of a slim disk.,2000) argued that the magnetic field is large in the NLS1s due to large pressure caused by the trapped photons in the inner region (radiation pressure dominated) of a slim disk.154 By requirnung Prag€Pu~Prada=aT?£3. one can estimate he temperature of the disk of the region that produced he flare.," By requiring $P_{mag}\le155P_{disc}\simeq P_{rad}=aT^3/3$, one can estimate the temperature of the disk of the region that produced the flare."156 This gives a νι23806V. which is somewhat ower than that of the lowest temmpcrature componcut in O'Brien et al," This gives a $kT_{disk}\ge 38 eV$, which is somewhat lower than that of the lowest temperature component in O'Brien et al."157 multiple black body model derived. from he broad band NAIAL spectimm.,'s multiple black body model derived from the broad band XMM spectrum.158 This perhaps explains why a rapid energetic flare could only observed in NLSIs with extremely steep soft X-ray., This perhaps explains why a rapid energetic flare could only observed in NLS1s with extremely steep soft X-ray.159 There is growing evidence for more massive black holes (BIT) iu radio-loud quasars than in radio quiet ones., There is growing evidence for more massive black holes (BH) in radio-loud quasars than in radio quiet ones.160 By assunune the lue cutting gas is virialized. Laor (2000) derived BID masses of the radio loud QSOs above 10° ML... significantly larger than in radio quiet QSOs.," By assuming the line emitting gas is virialized, Laor (2000) derived BH masses of the radio loud QSOs above $10^9$ $_\odot$, significantly larger than in radio quiet QSOs."161 This result is iu line with the receut identified correlation between the mass of the central DIT aud that of the bulge (Alagorrian et al., This result is in line with the recent identified correlation between the mass of the central BH and that of the bulge (Magorrian et al.162 1998. Cebhard et al.," 1998, Gebhard et al."163 2000. Ferrarese Meyitt 2000). together with the fact that the radio loud objects are hosted by elliptical galaxies.," 2000, Ferrarese Merritt 2000), together with the fact that the radio loud objects are hosted by elliptical galaxies."164 It is not uuderstood how the jet formation is related to the mass of the central DIT in AGN. the galactic DIT binaries also show super-huninal radio jets but their typical mass is around LOA...," It is not understood how the jet formation is related to the mass of the central BH in AGN, the galactic BH binaries also show super-luminal radio jets but their typical mass is around $_\odot$."165 Λοπιο et al. (, McLure et al. (1661999) found that both bhuuinous radio-loud ancl radio quiet quasars reside iu elliptical galaxies. sugeested massive DII for both type QSOs.,"1999) found that both luminous radio-loud and radio quiet quasars reside in elliptical galaxies, suggested massive BH for both type QSOs."167 Ou the other haud. it was suggested that the formation of jets night be related to the low accretion rate. such as in advection dominated accretion flow (c.g. Rees. Beechuan. Blandford Phinney 1982. Blandford Beechuan 1999) and/or with the spin of the DII ¢ Blauctord Zuajek 1977).," On the other hand, it was suggested that the formation of jets might be related to the low accretion rate, such as in advection dominated accretion flow (e.g, Rees, Begelman, Blandford Phinney 1982, Blandford Begelman 1999) and/or with the spin of the BH ( Blandford Znajek 1977)."168 It is worthy to mention that the radio ciaission from the BIT N-vay binarics also appear to be strong in the low state (Fender 2000). while it is suppressed in the high state.," It is worthy to mention that the radio emission from the BH X-ray binaries also appear to be strong in the low state (Fender 2000), while it is suppressed in the high state."169 However. evidence shows that the PISS 0558-501 possesses a low mass DIT and high accretion rate.," However, evidence shows that the PKS 0558-504 possesses a low mass BH and high accretion rate."170 The IL} width of this object is around 1250 kins 1 (Corbin 1997) and the optical contiuuun luminosity (1001}~2210% coves | from the V iaguitude and corrected for the Calactic reddening (Corbin Suuith 2000).," The $\beta$ width of this object is around 1250 km $^{-1}$ (Corbin 1997) and the optical continuum luminosity $\nu171L\nu(5100\AA) \simeq 2.2~10^{45}$ erg $^{-1}$ from the V magnitude and corrected for the Galactic reddening (Corbin Smith 2000)."172 By adopting the empirical Broad Line Reeion (BLR) size versus optical bhuninmositv relation. RoepezmLOA)/10ovess17 (aspi et al.," By adopting the empirical Broad Line Region (BLR) size versus optical luminosity relation, $R_{BLR}\simeq 18.65173[\lambda L\lambda(5100\AA)/10^{44} {\rm ergs~s}^{-1}]^{0.7}$ ( Kaspi et al."174 2000). audU.65|ALAG further assuming that the broad line region is virialized. we can derive a mass of central BIL of L5«107 solar mass.," 2000), and further assuming that the broad line region is virialized, we can derive a mass of central BH of $4.5\times 10^7$ solar mass."175 With this BID mass. the object is ciitting at super-Eddinetou luminosity.," With this BH mass, the object is emitting at super-Eddington luminosity."176" Note the mass of the central BI is far less than those seen in the radio loud quasars iu the PG sample (Laor 2000). which was interred with the same""C method."," Note the mass of the central BH is far less than those seen in the radio loud quasars in the PG sample (Laor 2000), which was inferred with the same method."177 We wish to poiut that the V-maenitude of PISS 0558-501 is LL97 aud its D-V. U-D values similar to those of 3C 273.," We wish to point that the V-magnitude of PKS 0558-504 is 14.97 and its B-V, U-B values similar to those of 3C 273."178 If it were in the north sky. it would be SCected as a PG QSO as well.," If it were in the north sky, it would be selected as a PG QSO as well."179 There are a umuber of uncertainties in the above estimation of the amass., There are a number of uncertainties in the above estimation of the mass.180 We have used the cmipirical Rerr relation., We have used the empirical $R_{BLR}-L_{opt}$ relation.181" The BER is photo-ionized bv arfextreme L,,;UY photons. the size may scale more adequately with the ionizing continuum than with the optical huuinositv."," The BLR is photo-ionized by far/extreme UV photons, the size may scale more adequately with the ionizing continuum than with the optical luminosity."182 Since PISS 0558-501 shows τον big due bunip as detected by NMM in optical-UV and. broac mand ταν (O'Brien et al., Since PKS 0558-504 shows very big blue bump as detected by XMM in optical-UV and broad band X-ray (O'Brien et al.183 2000). it has more ionizing Hux than a typical QSO with similar optical Nuuinosity. lus a lareer BLR.," 2000), it has more ionizing flux than a typical QSO with similar optical luminosity, thus a larger BLR."184 To eive the order of magnitude of lis correction. we estimate an optical to UV. spectra slope aze0.1. which is 0.1 lower than typical QSO value.," To give the order of magnitude of this correction, we estimate an optical to UV spectral slope $\alpha\simeq 0.1$, which is 0.4 lower than typical QSO value."185 If this spectrum extends to the Lyman lait. there is a actor of 2 more flux at this han for tvpica QSO spectrum.," If this spectrum extends to the Lyman limit, there is a factor of 2 more flux at this than for typical QSO spectrum."186 This will only increases the mass of DII oa factor of 1.5. which is still far less than the masses of DII« in radio loud PC QSOs.," This will only increases the mass of BH by a factor of 1.5, which is still far less than the masses of BHs in radio loud PG QSOs."187 Since by this wav the increase of ionizing photon flux also raises the bolometric uuimositv. this would uot lower the fraction accretion rate.," Since by this way, the increase of ionizing photon flux also raises the bolometric luminosity, this would not lower the fraction accretion rate."188 Another coucern is that the BLR is in a flat structure. so iat the velocity depeuds strongly on the inclination of je system.," Another concern is that the BLR is in a flat structure, so that the velocity depends strongly on the inclination of the system."189 There is evidence for the auisotropx of BLR velocity in radio loud ACNs. aud the liue width is mud to be auti-correlated with the core dominance of je radio source.," There is evidence for the anisotropy of BLR velocity in radio loud AGNs, and the $\beta$ line width is found to be anti-correlated with the core dominance of the radio source."190 There is no such radio data for the PISs 3558-50L. so it remains possible that we might see a fat BLR from top in this object aud the mass of ceutral DIT was under-estimated.," There is no such radio data for the PKS 0558-504, so it remains possible that we might see a flat BLR from top in this object and the mass of central BH was under-estimated."191 However. the core dominated radio source uxuallv display flat spectra. instead of the steep radio spectrin in these objects (see 1.11.," However, the core dominated radio source usually display flat spectrum, instead of the steep radio spectrum in these objects (see 4.1)."192 Teh resolution radio observation. however. is necessary to address this.," High resolution radio observation, however, is necessary to address this."193 The other two radio loud NLSIs (ROB JOOLL)193 aud RN JOLS12-1258) have even lower DIT masses aud higher accretion rates;, The other two radio loud NLS1s (RGB J0044+193 and RX J0134.2-4258) have even lower BH masses and higher accretion rates.194 Using the same method. we fiud that the BID masses are 1.6«10° and 10* AL. for RGB JO0111193 and RX J0131.2-1258.. respectively.," Using the same method, we find that the BH masses are $1.6\times 10^7$ and $10^7$ $_\odot$ for RGB J0044+193 and RX J0134.2-4258, respectively."195 These two objects show the same characteristics iu the SED as PISS 0555-501: a very big blue bump imdicated by its fat optical to UV continuum aud huge κο X-ray excess. steep radio," These two objects show the same characteristics in the SED as PKS 0558-504: a very big blue bump indicated by its flat optical to UV continuum and huge soft X-ray excess, steep radio"196What could cause different time delavs at. cülferent epochs?,What could cause different time delays at different epochs?197 There are two obvious possibilities — clilferent physical size scales at. different epochs. or changes in the Doppler factor of the jet. (since. the. celay between different frequencies originates in the relativistically [owing meclium).," There are two obvious possibilities – different physical size scales at different epochs, or changes in the Doppler factor of the jet (since the delay between different frequencies originates in the relativistically flowing medium)."198 In reality there may be some combination of these. and other. effects.," In reality there may be some combination of these, and other, effects."199 If the first explanation is the main reason however. we may expect some observational relation. between radio ux evels and delays.," If the first explanation is the main reason however, we may expect some observational relation between radio flux levels and delays."200 Phe simultaneous 8 15 Gllz light curves oesented in Pooley Fender (1997) had lux densities in he range 50100 mJy., The simultaneous 8 15 GHz light curves presented in Pooley Fender (1997) had flux densities in the range 50–100 mJy.201 Phe oscillations presented. in this paper have a range of about 40.100 mv., The oscillations presented in this paper have a range of about 40–100 mJy.202 The amplitudes reported in. Mirabel et al. (, The amplitudes reported in Mirabel et al. (2031998). up to 100 my. appear os somewhat. larger.,"1998), up to 100 mJy, appear to be somewhat larger."204 Thus the longer delays may. be due to jet structure which is larger overall. ancl we would expec eiven more data. to see a correlation between delay. length— and oscillation amplitude.," Thus the longer delays may be due to jet structure which is larger overall, and we would expect, given more data, to see a correlation between delay length and oscillation amplitude."205 Η the explanation is instead due to a changing Doppler factor. this may be most naturally explained as being due to a change in the angle of the jet to the line of sight (although we cannot rule out changes in the bulk Lorentz factor). for example if the jet is precessing.," If the explanation is instead due to a changing Doppler factor, this may be most naturally explained as being due to a change in the angle of the jet to the line of sight (although we cannot rule out changes in the bulk Lorentz factor), for example if the jet is precessing."206 For a bulk Lorentz factor of 5 and angle of jet to the line of sight of ~65 degrees. as calculated. for GRS 1915]105 ejections for a distance of 11 kpc (render et al.," For a bulk Lorentz factor of 5 and angle of jet to the line of sight of $\sim 65$ degrees, as calculated for GRS 1915+105 ejections for a distance of 11 kpc (Fender et al."207 1999). a swing of z20 degrees would be required in order to change the Doppler factor by the factor of ~2 indicated hy table 1 (for larger Lorentz factors a smaller swing in angle is required. but this would have further implications for e.g. energeties).," 1999), a swing of $\ga 20$ degrees would be required in order to change the Doppler factor by the factor of $\sim 2$ indicated by table 1 (for larger Lorentz factors a smaller swing in angle is required, but this would have further implications for e.g. energetics)."208" In the event that varving delavs were caused by a periodically: varving jet angle to the line of sight. we might expect to see. in the long-term. a repeating pattern of ‘long and ""short delays."," In the event that varying delays were caused by a periodically varying jet angle to the line of sight, we might expect to see, in the long-term, a repeating pattern of `long' and `short' delays."209 1n a set of high-sensitivity cdual-cfrequeney racio observations of GRS 19151105. with ΑΟΔΑ during. a period of radio oscillations we have captured. in. detail some kev “haracteristics., In a set of high-sensitivity dual-frequency radio observations of GRS 1915+105 with ATCA during a period of radio oscillations we have captured in detail some key characteristics.210 Firstly we have clearly shown that the Gllz ractio μα»eetrum. of CRS 1915]105. repeatedly: eveles between pticallv thick and thin during periods of radio oscillations., Firstly we have clearly shown that the GHz radio spectrum of GRS 1915+105 repeatedly cycles between optically thick and thin during periods of radio oscillations.211 While this could be qualitatively appreciated from: previous μαxwse observations and simple concepts. we have for the first time captured the strong quasi-periodic signal in the radio spectrum.," While this could be qualitatively appreciated from previous sparse observations and simple concepts, we have for the first time captured the strong quasi-periodic signal in the radio spectrum."212 Lt can now be directly. appreciated. from Fig 2. that were we to average over multiple. oscillation events (for e.g. a weaker source or a shorter oscillation period). we would measure a spectral index which was a combination of both optically thick and optically thin emission. with an average close to zero. as seen in hard states of black hole candidate X-ray binaries (Fender 2001).," It can now be directly appreciated, from Fig 2, that were we to average over multiple oscillation events (for e.g. a weaker source or a shorter oscillation period), we would measure a spectral index which was a combination of both optically thick and optically thin emission, with an average close to zero, as seen in hard states of black hole candidate X-ray binaries (Fender 2001)."213 The transition from. optically thin to flat/inverted. racio spectra following the outburst is also reminiscent of the iud state transients (see e.g. Fig 3 of Fender 2001)., The transition from optically thin to flat/inverted radio spectra following the outburst is also reminiscent of the hard state transients (see e.g. Fig 3 of Fender 2001).214 In addition. the ~I level of linear polarisation associated with the oscillations at both frequencies is comparable to hat measured for the Dat spectrum in the hard states of wo black hole candidates. CS 2033|338/V404. Cve (Llan LHLjellming 1992). and GX 339-4 (Corbel et al.," In addition, the $\sim2151$ level of linear polarisation associated with the oscillations at both frequencies is comparable to that measured for the flat spectrum in the hard states of two black hole candidates, GS 2033+338/V404 Cyg (Han Hjellming 1992), and GX 339-4 (Corbel et al."216 2000). and his level of linear polarisation may be a generic feature of such flat spectra in DIC. low/hard states (Fender 2001).," 2000), and this level of linear polarisation may be a generic feature of such flat spectra in BHC low/hard states (Fender 2001)."217 The similarity of the linear polarisation. position angle for hh the optically thin [lare and. oscillations may indicate a preferred orientation of magnetic field for the outllow., The similarity of the linear polarisation position angle for both the optically thin flare and oscillations may indicate a preferred orientation of magnetic field for the outflow.218 In addition in this case we have measured the spectrum of the inear polarisation. and find it also to be approximately Lat.," In addition in this case we have measured the spectrum of the linear polarisation, and find it also to be approximately flat."219 Finally. we have performed a cross-correlation analysis of the two radio lighteurves in order to establish a mean cdelav.," Finally, we have performed a cross-correlation analysis of the two radio lightcurves in order to establish a mean delay."220 This is the most thorough analysis of the time dclays to date: previous attempts were by eve! or involved. only the peaks of single. events., This is the most thorough analysis of the time delays to date; previous attempts were `by eye' or involved only the peaks of single events.221 Comparing all the reported frequeney-dependent time delays for the oscillations from the literature. we find that there are significant dilferences at dillerent epochs.," Comparing all the reported frequency-dependent time delays for the oscillations from the literature, we find that there are significant differences at different epochs."222 The longest time delays reported seem to be associated with the large-amplitude: oscillations in Mirabel et al. (, The longest time delays reported seem to be associated with the large-amplitude oscillations in Mirabel et al. (2231998). which max indicate a relation between event amplitude and subsequent time delay.,"1998), which may indicate a relation between event amplitude and subsequent time delay."224 However. at this stage a changing Doppler factor or some other cllect cannot be ruled out.," However, at this stage a changing Doppler factor or some other effect cannot be ruled out."225 Detailed. observations such as these are crucial to a quantitative understanding of the physies involved. in jet, Detailed observations such as these are crucial to a quantitative understanding of the physics involved in jet226 (B=105 idensities. B=10!1—107 as , $B \gtrsim 10^{13}$ $B \approx 10^{14}-10^{15}$ 227"eas found in the broad A0,512.8 (3.3) transition. which bears a close resemblance to the CO peaks.","is found in the broad $\Delta\,v_{1/2}$ (3,3) transition, which bears a close resemblance to the CO peaks."228 The velocity dispersion. image of the (3.38) ine in Figure 9 shows clearly that the cloud disruption originates from the western or W28 SNIU side. supporting he results of Arikawactal.(1999) who showed that the woad “CO ((32) emission is found. preferentially towards he WS side in contrast to 17CO ((1:O) which extends racially further away.," The velocity dispersion image of the (3,3) line in Figure \ref{fig:veldisp} shows clearly that the cloud disruption originates from the western or W28 SNR side, supporting the results of \citet{arikawa} who showed that the broad $^{12}$ (3–2) emission is found preferentially towards the W28 side in contrast to $^{12}$ (1–0) which extends radially further away."229 Lhe Nanten2 CO (21) emission ikelv traces a mixture of shocked ancl unshocked gas ancl detailed: velocity. dispersion studies of this (21) emission are currently underway., The Nanten2 $^{12}$ (2–1) emission likely traces a mixture of shocked and unshocked gas and detailed velocity dispersion studies of this (2–1) emission are currently underway.230 The lack of strong broad-band Lh eatures towards Core 2 (Figure 7)) suggests that the NIL; excitation is not due to star formation processes. although some fraction of the weak LR emission seen here is due o shocked. ο (Neufeld.ctal.2007:Alarquez-Lugo 2010).," The lack of strong broad-band IR features towards Core 2 (Figure \ref{fig:IRdata}) ) suggests that the $_{3}$ excitation is not due to star formation processes, although some fraction of the weak IR emission seen here is due to shocked $_2$ \citep{neufeld,marquez-lugo}."231". We also note there is a weak LIC'N feature seen in 10 deep pointing which is an indicator of hot. gas-phase ""hemistry.", We also note there is a weak $_3$ N feature seen in the deep pointing which is an indicator of hot gas-phase chemistry.232 Lt is also quite striking that a grouping of the OLLI masers appear to surround the region where the (3.3) emission is most disrupted. (Figure 9)). racially away [rom 1 Wes SNR.," It is also quite striking that a grouping of the OH masers appear to surround the region where the (3,3) emission is most disrupted (Figure \ref{fig:veldisp}) ), radially away from the W28 SNR."233 This would be further evidence in support of 1 ΑΝ SNIt as the source of molecular cloud. disruption. ux overall would tend o cisfavour physical inlluence from 1f neighbouring SNR GG6.67-0.42. which at. present has an unknown distance.," This would be further evidence in support of the W28 SNR as the source of molecular cloud disruption, and overall would tend to disfavour physical influence from the neighbouring SNR G6.67-0.42, which at present has an unknown distance."234 Phe broad molecular line. regions discussed. by Reachetal.(2005) also generally cluster towards the broadest NIL; emission (see Figure 9))., The broad molecular line regions discussed by \citet{reach} also generally cluster towards the broadest $_3$ emission (see Figure \ref{fig:veldisp}) ).235 The broadening of the NIL; (03.3). (6.6) anc probably (9.9) lines are dominated by non-thermal component(s) and we can estimate the additional kinetic energy. VM reeuired to achieve this from: whereAZ is the mass of the broac-line gas and Neg is the FWHAL !] of the line due to additional non-hermal kinetic processes.," The broadening of the $_3$ (3,3), (6,6) and probably (9,9) lines are dominated by non-thermal component(s) and we can estimate the additional kinetic energy $W_{\rm kin}$ required to achieve this from: where$M$ is the mass of the broad-line gas and $\Delta v_{\rm kin}$ is the FWHM $^{-1}$ ] of the line due to additional non-thermal kinetic processes."236" Using the non-thermal line width Ain=7.5 and mass lower limit Af=1300 MM. [rom our radiative ransfer modeling in refssecirtmC'2 of mapping-averaged NIL, spectra we herefore calculate Wye0.7Loir cere."," Using the non-thermal line width $\Delta v_{\rm kin}=7.5$ and mass lower limit $M=1300$ $_\odot$ from our radiative transfer modeling in \\ref{ssec:rtmC2}237 of mapping-averaged $_3$ spectra we therefore calculate $_{KE}>0.7\times10^{48}$ erg."238 pjThis energy ower limit is within a [actor few of the kinetic energy (~3.107 cere) deposited into the MAL. of gas traced w shocked ((3.2) from Arikawactal.1999..," This energy lower limit is within a factor few of the kinetic energy $\sim\,3\times10^{48}$ erg) deposited into the $_\odot$ of gas traced by shocked $^{12}$ (3–2) from \citealt{arikawa}."239 Over the 0 to range. for which the Nanten CO ((10) emission shows excellent overlap with the TeV eanma-ray source LESS 1501-9399. the mass of the NE molecular cloud is ~2107 MM...," Over the 0 to range, for which the Nanten $^{12}$ (1–0) emission shows excellent overlap with the TeV gamma-ray source HESS J1801-233, the mass of the NE molecular cloud is $\sim 2\times 10^{4}$ $_{\odot}$ ."240 Over a wider 0 to, Over a wider 0 to241Iuminous infrared galaxy 2220 (1140.vanderTaketal.2005).,luminous infrared galaxy 220 \citep[$\rm H_3O^+$][]{VdTak08}.242. Therefore. 882 is still outstanding not onlv as a PDR dominated galaxy. but by the underabundance of complex molecules such as CHOI. LINCOor SiO (MauersbergerMartínetal. 2006a.b).. evidence for the lack of large amounts of dense molecular material which would potentially fuel its nuclear starburst as compared to other starburst galaxies like 2253 (Martínetal.2009).," Therefore, 82 is still outstanding not only as a PDR dominated galaxy, but by the underabundance of complex molecules such as $_3$ OH, HNCOor SiO \citep{Mauers93,Martin06a,Martin06b}, evidence for the lack of large amounts of dense molecular material which would potentially fuel its nuclear starburst as compared to other starburst galaxies like 253 \citep{Martin09}."243.. Our data in combination with the ΗΝο abundances (Martinetal.2009) indicate that the moleeular clouds in 882 are different [rom those in NGC253., Our data in combination with the HNCO abundances \citep{Martin09} indicate that the molecular clouds in 82 are different from those in NGC253.244 Although having a similar overall PDR. component. the clouds in 2253 have to be more massive ancl have larger column densities those in 882.," Although having a similar overall PDR component, the clouds in 253 have to be more massive and have larger column densities those in 82."245" This work has been partially supported by (he Spanish Ministerio de Ciencia e Innovacionn under project ESP2001-65312-CO02-01. and by the ""Comunidad cle Madrid Government under PRICIT project $-0505/ESP-0237 CASTROCAMD. JCAIT.."," This work has been partially supported by the Spanish Ministerio de Ciencia e Innovaciónn under project ESP2007-65812-C02-01, and by the “Comunidad de Madrid” Government under PRICIT project S-0505/ESP-0237 (ASTROCAM). ."2462002).,.247. A nageing deficiency of p-process production in the mass range A = 92 - 124 still persisted though., A nagging deficiency of $p$ -process production in the mass range A = 92 - 124 still persisted though.248" The production of °°""! Mo posed a particular problem since. unless (he star had previously experienced a strong s-process. enhancing the abundance of seed above A = 95. there simply was not enough seed."," The production of $^{92,94}$ Mo posed a particular problem since, unless the star had previously experienced a strong $s$ -process, enhancing the abundance of seed above A = 95, there simply was not enough seed."249 In massive stars (he s-process does not go above mass 90 and so (he necessary seed enhancement does not occur., In massive stars the $s$ -process does not go above mass 90 and so the necessary seed enhancement does not occur.250" Hoffmanetal.(1996). found that large abundances of some p-nuclei. and ""Mo in particular. could be synthesized in the neutrino-powered wind blowing from a voung neutron star (seealsoDuncan.Shapiro.&Wasserman 1956)."," \citet{hof96} found that large abundances of some $p$ -nuclei, and $^{92}$ Mo in particular, could be synthesized in the neutrino-powered wind blowing from a young neutron star \citep[see251also][]{dun86}."252". While this wind had chielly been seen as a wav of making the rprocess (Wooslevetal. 1994).. for electron mole numbers. Y;z0.485. the p-nuclei ""Zn. Se. “Ie. ! Sr. and °°? Mo were produced in great abundance."," While this wind had chiefly been seen as a way of making the $r$ -process \citep{woo94}, , for electron mole numbers, $Y_e253\approx 0.485$, the $p$ -nuclei $^{64}$ Zn, $^{74}$ Se, $^{78}$ Kr, $^{84}$ Sr, and $^{92}$ Mo were produced in great abundance."254 It is important (o note in this regard (hat. while Y.=0.485 is nominally neutron rich (Y; = ().5 corresponds (o neutron. proton equality). it is still a lot more proton-rich (han the p- themselves (Z/N [or  Mo = 0.457). so the nucleonic gas contained some [ree protons.," It is important to note in this regard that, while $Y_e = 0.485$ is nominally neutron rich $Y_e$ = 0.5 corresponds to neutron, proton equality), it is still a lot more proton-rich than the $p$ -nuclei themselves (Z/N for $^{92}$ Mo = 0.457), so the nucleonic gas contained some free protons."255 The p-nuclei here were also primary. in (he sense (hat a star with no initial metallicity would still make (he same composition in ils neutrino wind.," The $p$ -nuclei here were also primary, in the sense that a star with no initial metallicity would still make the same composition in its neutrino wind."256" There were potential problems. however. in that the ejection of only asmall amount of mass with Y. just a little lower than 0.485 resulted in disastrous overproduction of N = 50 nuclei like “Sr. Y. and ""Zr."," There were potential problems, however, in that the ejection of only asmall amount of mass with $Y_e$ just a little lower than 0.485 resulted in disastrous overproduction of N = 50 nuclei like $^{88}$ Sr, $^{89}$ Y, and $^{90}$ Zr."257 Also the neutron-rich wind failed to produce acequate amounts of p-process nuclei above A = 92., Also the neutron-rich wind failed to produce adequate amounts of p-process nuclei above A = 92.258 Though this paper focuses on early proton-rich outflows. the SN model we study is calculated. (ο eject a sizable amount of neutron-rich material.," Though this paper focuses on early proton-rich outflows, the SN model we study is calculated to eject a sizable amount of neutron-rich material."259 Ho remains to be seen if verv recent. simulations predict neutron-rich outflows that satishy the conditions needed [or efficient svnthesis of °?Mo. or if neutrino interactions facilitate production of °°\lo in (he neutron rich ejecta predicted by these models (Fuller&Alever1995).," It remains to be seen if very recent simulations predict neutron-rich outflows that satisfy the conditions needed for efficient synthesis of $^{92}$ Mo, or if neutrino interactions facilitate production of $^{92}{\rm Mo}$ in the neutron rich ejecta predicted by these models \citep{ful95}."260. Based upon calculations bx Jim Wilson. Qian&Wooslev(1996) pointed out. that yo in the wind would naturally evolve though the points necessary to make these p-nuclei and would actually start with a value greater than 0.5.," Based upon calculations by Jim Wilson, \citet{qia96} pointed out that $Y_e$ in the wind would naturally evolve though the points necessary to make these $p$ -nuclei and would actually start with a value greater than 0.5."261 As other detailed models for core- supernovae became available. nucleosyntliesis was explored in (his “hot. proton-rich bubble” bv Pruetetal.(2005).. Frohlichetal. (2004).. ancl Frohlichetal. (2005)..," As other detailed models for core-collapse supernovae became available, nucleosynthesis was explored in this “hot, proton-rich bubble” by \citet{pru05}, \citet{fro04}, and \citet{fro05}. ."262 The latter two stuclies found. substantial production of nuclei up to A = 84. including some nuclei (racitionally attributed to the p-process.," The latter two studies found substantial production of nuclei up to A = 84, including some nuclei traditionally attributed to the $p$ -process."263" It seems probable that these winds have also contributed appreciably to the solar abundances of PSc.. Ti and ""! Zn. and. possibly in a measurable wav. to other rare abundances in metal poor stars."," It seems probable that these winds have also contributed appreciably to the solar abundances of $^{45}$ Sc, $^{49}$ Ti and $^{64}$ Zn, and, possibly in a measurable way, to other rare abundances in metal poor stars."264 However. since (hesesime nuclei were already madeby other. processes (Wooslev.Heger.&Weaver2002: 2002).. there seemed to be no clear ciagnosticof the proton-rich wind.," However, since thesesame nuclei were already madeby other processes \citep{Woo02,Rau02}, , there seemed to be no clear diagnosticof the proton-rich wind."265 llere. following the suggestion of Frohlichetal. (2005).. we have revisited our caleulations," Here, following the suggestion of \citet{fro05}, , we have revisited our calculations"266A plot of the gaseous mass density versus the line-of-sight CO velocity dispersion for the five ULICs listed iu Table 1.,A plot of the gaseous mass density versus the line-of-sight CO velocity dispersion for the five ULIGs listed in Table 1.267 The CO velocity dispersion. 6. has been estimated bv treating each CO line profile as Caussian.," The CO velocity dispersion, $\sigma$, has been estimated by treating each CO line profile as Gaussian."268 Thus. 6=Avpgwipnr/2.351.," Thus, $\sigma = \Delta {\rm v_{FWHM}} /2692.354$."270 The small and laree circles connected by lines represent mass density ranges limited using CO size estimates 1 and 2. respectively (see 85).," The small and large circles connected by lines represent mass density ranges limited using CO size estimates 1 and 2, respectively (see 5)."271 Also plotted are the stellar mass densities and line-ofsight velocity dispersions of elliptical galaxies with both power-law lisht profiles aud cores., Also plotted are the stellar mass densities and line-of-sight velocity dispersions of elliptical galaxies with both power-law light profiles and cores.272 The latter data are taken from Faber et al. (, The latter data are taken from Faber et al. (2731997).,1997).274"calculating the principal axes e. b. and e (ee.28?)., where e. b. aud e denote the major. intermeciate. aud ninor axis. respectively.","calculating the principal axes $a$ , $b$ , and $c$ \citep[e.g.,][]{Springel_etal04,275 Hayashi_etal07,Abadi_etal10}, where $a$, $b$, and $c$ denote the major, intermediate, and minor axis, respectively."276 The trasgal DAL halo models are constructed via successive qguereers 7., The triaxial DM halo models are constructed via successive mergers \citet{Moore_etal04}.277" Tho initial system is a το, 000-particle. isotropic. spherically-svuuuetric— model eenecrated by sampling the ? distribution function (ην."," The initial system is a $75,000$ -particle, isotropic, spherically-symmetric model generated by sampling the \citet{Hernquist90} distribution function \citep{Kazantzidis_etal04b}."278 À nearly prolate recunant results from a head-ou collision between two such Heruquist. models: more riaxial configurations are built by mereing this reiinaut with other Iheruquist models ou parabolic orbits of various inclinations relative to the principal axes of the first reninant., A nearly prolate remnant results from a head-on collision between two such Hernquist models; more triaxial configurations are built by merging this remnant with other Hernquist models on parabolic orbits of various inclinations relative to the principal axes of the first remnant.279 Each final triaxial halo consists of 300.000 articles.," Each final triaxial halo consists of $300,000$ particles."280 After the mergers are complete. we evolve tle relunauts for several dvuamiucal times in order to ensure hat equilibrium is achieved.," After the mergers are complete, we evolve the remnants for several dynamical times in order to ensure that equilibrium is achieved."281 The final coufiguratious lave nass distributions that can be well approximated wo NEW profiles (7)| Gnu the regions relevant to the yresent study} and are thus cousistent with the results of costnological CDM. simulations., The final configurations have mass distributions that can be well approximated by NFW profiles \citep{Navarro_etal96} (in the regions relevant to the present study) and are thus consistent with the results of cosmological CDM simulations.282 We adopt two different halo models in order to grow central disks., We adopt two different halo models in order to grow central disks.283 The first Guodel A) is a nearly prolate πιο renminant. with trianiality pariancter (?) Fo—(a?Pyf(a*ο}~(UN for most radii.," The first (model A) is a nearly prolate merger remnant, with triaxiality parameter \citep{Franx_etal91} $T = (a^2 - b^2) / (a^2 - c^2) \sim2840.8$ for most radii."285 The second (model B) is stronely triaxial Z7—(0.1 in the iuner reelons. increasing to £~0.6 in the outskirts.," The second (model B) is strongly triaxial; $T286\sim 0.4$ in the inner regions, increasing to $T \sim 0.6$ in the outskirts."287 These shapes are fairly typical of those fouud in cosmological simulations (7)., These shapes are fairly typical of those found in cosmological simulations \citep{Hayashi_etal07}.288. The spin paramcter A of halos A and D are 0.021 and 0.040. respectively. which are typical of cosinological halos (?)..," The spin parameter $\lambda$ of halos A and B are 0.024 and 0.040, respectively, which are typical of cosmological halos \citep{Maccio_etal08}."289" The circular velocity aud axial ratio profiles of these halos are preseuted in Figures 1 and 2.. respectively,"," The circular velocity and axial ratio profiles of these halos are presented in Figures \ref{fig1} and \ref{fig2}, respectively."290 We scale the parameters of halos A aud D to match the 1nüass and concentration of CDM. halos with maxima cireular velocity Vyas~120ας|.," We scale the parameters of halos A and B to match the mass and concentration of CDM halos with maximum circular velocity $V_{\rm max}291\sim 120 \kms$."292 For a halo of virial mass Ma=5«101AZ. and concentration e~9.3 (2). the circular velocityprofile peaks at riascLS spe.," For a halo of virial mass $M_{\rm vir} = 5 \times29310^{11} M_{\odot}$ and concentration $c \sim 9.3$ \citep{Neto_etal07}, the circular velocityprofile peaks at $r_{\rm max} \sim 48$ kpc."294 Disk galaxies forming in our triaxial halos mst have rotation speeds comparableto Vyas du order to satisfy simuultaneouslv the uormalization of the Tully-Fisher, Disk galaxies forming in our triaxial halos must have rotation speeds comparableto $V_{\rm max}$ in order to satisfy simultaneously the normalization of the Tully-Fisher295"In this work, we carried out collapse simulations of magnetised and rotating molecular cloud cores using the Smoothed Particle Magnetodynamics (SPMHD) method.","In this work, we carried out collapse simulations of magnetised and rotating molecular cloud cores using the Smoothed Particle Magnetodynamics (SPMHD) method."296" We reproduced important features of protostellar outflows, qualitatively in good agreement with results reported in the literature."," We reproduced important features of protostellar outflows, qualitatively in good agreement with results reported in the literature."297" Furthermore, we computed the ratio Btor/Bpoi and thus give a quantitative proof of the evolution of the toroidal part of the magnetic field connected with the outflow."," Furthermore, we computed the ratio $B_{\textrm{tor}}/B_{\textrm{pol}}$ and thus give a quantitative proof of the evolution of the toroidal part of the magnetic field connected with the outflow."298 Rendered plots were made using the software written by Daniel Price (?).., Rendered plots were made using the software written by Daniel Price \citep{Price2007hc}.299 aacknowledges granting of computer time from John von Neumann-Institute for, acknowledges granting of computer time from John von Neumann-Institute for300there are already papers based on high resolution spectra (see next Section). but we improve on past studies both on statistics (for most of them only two stars were analyzed) and S/N ratio.,"there are already papers based on high resolution spectra (see next Section), but we improve on past studies both on statistics (for most of them only two stars were analyzed) and S/N ratio."301 The paper is organized as follows: in Sect., The paper is organized as follows: in Sect.302 2 we describe the observational samples. data reduction. and abundance analysis: the results are reported in Sect.," 2 we describe the observational samples, data reduction, and abundance analysis; the results are reported in Sect."303 3 and discussed in Sect., 3 and discussed in Sect.304 4: a summary and conclusion are given in Sect., 4; a summary and conclusion are given in Sect.305 5., 5.306 This is the most distant cluster in our sample (actually. the most distant known open cluster) and one of the most metal poor ones.," This is the most distant cluster in our sample (actually, the most distant known open cluster) and one of the most metal poor ones."307 The first photometric study of this cluster is the CCD BVI survey by Kaluzny (1994)) who estimated an age of «4 Gyr and a distance of 10.5 kpe., The first photometric study of this cluster is the CCD $BVI$ survey by Kaluzny \cite{kalBe29}) ) who estimated an age of $\sim$ 4 Gyr and a distance of 10.5 kpc.308 Tosi et al. (2004)), Tosi et al. \cite{tosi04}) )309 in a CCD BVT study found. depending on the adopted stellar models. an age of 3.4 or 3.7 Gyr. Gu— M)g= 15.6 or 15.8 (implying Ay. = 21 or 22 κρο). with E(B—V) = 0.13 or 0.10. and metallicity lower than solar (Z2 0.006 or 0.004). depending on the adopted stellar models.," in a CCD $BVI$ study found, depending on the adopted stellar models, an age of 3.4 or 3.7 Gyr, $(m-M)_0$ = 15.6 or 15.8 (implying $R_{\rm gc}$ = 21 or 22 kpc), with $E(B-V)$ = 0.13 or 0.10, and metallicity lower than solar (Z= 0.006 or 0.004), depending on the adopted stellar models."310 These values were confirmed by Bragaglia. Held. Tosi (2005)). who presented radial velocity measurements for red giant branch (RGB) stars. and by Bragaglia Tosi (2006)) who give the followingparameters based on the Padova evolutionary tracks: age ~ 3.7 Gyr. Gi— M)oz15.6. ECB— V)=0.12. Zz0.004. distance 13.05 kpe.," These values were confirmed by Bragaglia, Held, Tosi \cite{bragaglia05}) ), who presented radial velocity measurements for red giant branch (RGB) stars, and by Bragaglia Tosi \cite{bt06}) ) who give the followingparameters based on the Padova evolutionary tracks: age $\sim$ 3.7 Gyr, $(m-M)_{0}$ =15.6, $E(B-V)$ =0.12, $Z$ =0.004, distance 13.05 kpc."311 Spectroscopic investigations of Be 29 were presented by Carraro et al. (2004)).," Spectroscopic investigations of Be 29 were presented by Carraro et al. \cite{carraro04}) ),"312 Yong et al. (2005)).," Yong et al. \cite{yong05}) ),"313 and Frinchaboy et al. (2006))., and Frinchaboy et al. \cite{frincha06}) ).314 Carraro et al., Carraro et al.315 and Yong et al., and Yong et al.316 derived a metallicity [Fe/H|2-0.44 (for two red clump stars) and —0.54 (for two stars near the RGB tip). respectively.," derived a metallicity $-0.44$ (for two red clump stars) and $-0.54$ (for two stars near the RGB tip), respectively."317 They also presented results for abundances of several elements in addition to Fe., They also presented results for abundances of several elements in addition to Fe.318 Frinchaboy et al., Frinchaboy et al.319 investigated the radial velocity of the cluster. concluding that its properties are in agreement with the Galactic anticenter stellar structure. also known as the Monoceros Ring.," investigated the radial velocity of the cluster, concluding that its properties are in agreement with the Galactic anticenter stellar structure, also known as the Monoceros Ring."320" Also Be 20 is a very distant cluster. with a distance of 8.4 kpe from the Sun and A, of about 16 kpe."," Also Be 20 is a very distant cluster, with a distance of 8.4 kpc from the Sun and $R_{\rm gc}$ of about 16 kpc."321 The first photometric study of Be 20 1s the CCD V7 survey by MacMinn et al. (1994)):, The first photometric study of Be 20 is the CCD $VI$ survey by MacMinn et al. \cite{mcminn}) );322 they estimated an age of «5-6 Gyr. a reddening E(V—-1)=0.16. a metallicity 20.23 and 640—M)é 2105.0; their data were used for the present paper.," they estimated an age of $\sim$ 5–6 Gyr, a reddening $E(V-I)$ =0.16, a metallicity $-$ 0.23 and $(m-M)_{V}$ =15.0; their data were used for the present paper."323 Recently Andreuzzi et al. (2007.. 2008))," Recently Andreuzzi et al. \cite{abt07}, \cite{abt08}) )"324 presented a new photometry. obtaining results in very good agreement with MacMinn et al. (1994)).," presented a new photometry, obtaining results in very good agreement with MacMinn et al. \cite{mcminn}) )."325 Two spectroscopic investigations of Be 20 were carriedout so far., Two spectroscopic investigations of Be 20 were carriedout so far.326 Friel et al. (2002)).," Friel et al. \cite{friel02}) ),"327 on the basis of low resolution spectra of 6 stars. derived an average [Fe/H] value of -0.61+0.14 dex.," on the basis of low resolution spectra of 6 stars, derived an average [Fe/H] value of $-0.61\pm 0.14$ dex."328 Yong et al. (2005)), Yong et al. \cite{yong05}) )329 derived [Fe/H|2—0.45 and -0.53 for two stars near the RGB tip observed at high resolution., derived $-$ 0.45 and $-$ 0.53 for two stars near the RGB tip observed at high resolution.330 This 15 one of the oldest open clusters in the Galaxy. with an age in the range 6-10 Gyr (but the upper value is most probably spurious. see the discussion in Bragaglia Tost 2006)).," This is one of the oldest open clusters in the Galaxy, with an age in the range 6–10 Gyr (but the upper value is most probably spurious, see the discussion in Bragaglia Tosi \cite{bt06}) )."331 It is also one of the few known old open clusters inside the Solar circle (δω~7.5 kpc)., It is also one of the few known old open clusters inside the Solar circle $R_{\rm gc}\sim7.5$ kpc).332 Photometric studies of the cluster were carried out by Janes Phelps (1994)). Mazur. Krzeminski. Kaluzny (1905). and Gozzoli et al. (1996));," Photometric studies of the cluster were carried out by Janes Phelps \cite{JP94}) ), Mazur, Krzeminski, Kaluzny \cite{mazur95}) ), and Gozzoli et al. \cite{gozzoli}) );"333 not surprisingly. given the low Galactic latitude and the direction toward the center. they all agree on a high reddening values. ranging from E(B—V)=0.22 up to 0.34 (depending on the assumed metallicity).," not surprisingly, given the low Galactic latitude and the direction toward the center, they all agree on a high reddening values, ranging from $E(B-V)$ =0.22 up to 0.34 (depending on the assumed metallicity)."334 The Gozzoli et al., The Gozzoli et al.335 results have been recently revised by Bragaglia Tosi (2006)). who find a best fit solution for solar metallicity. distance modulus Gv—M)5212.2. E(B—V)=0.30. and age 6 Gyr.," results have been recently revised by Bragaglia Tosi \cite{bt06}) ), who find a best fit solution for solar metallicity, distance modulus $(m-M)_{0}$ =12.2, $E(B-V)$ =0.30, and age 6 Gyr."336 The metallicity of Cr 261 has been measured by means of low resolution spectroscopy by Friel et al. (2002)):, The metallicity of Cr 261 has been measured by means of low resolution spectroscopy by Friel et al. \cite{friel02}) ):337 from the analysis of 21 giants they derived [Fe/H] = —0.1620.13., from the analysis of 21 giants they derived [Fe/H] = $-$ $\pm$ 0.13.338 In a more recent study the same group determined the metallicity from high resolution data of four giants. obtaining [Fe/H] = —0.22+0.05 (Friel et al. 2003)).," In a more recent study the same group determined the metallicity from high resolution data of four giants, obtaining [Fe/H] = $-$ $\pm$ 0.05 (Friel et al. \cite{friel03}) )."339 A higher metal content is reported by Carretta et al. (2005)).," A higher metal content is reported by Carretta et al. \cite{carretta05}) ),"340 who derived [Fe/H] = —0.0340.03 (1.e. solar metallicity) and proposed as probable reasons for the different result differences in the model atmospheres. in the spectral resolution and equivalent width (EW) measurements. and in the atomic parameters (oscillator strength).," who derived [Fe/H] = $-$ $\pm$ 0.03 (i.e. solar metallicity) and proposed as probable reasons for the different result differences in the model atmospheres, in the spectral resolution and equivalent width $EW$ ) measurements, and in the atomic parameters (oscillator strength)."341 Deep CCD photometry for Mel 66 was provided by Kassis et al. (1997)).," Deep CCD photometry for Mel 66 was provided by Kassis et al. \cite{kassis}) ),"342 who quote an age of 441 Gyr and ο— M)o=13.2. with [Fe/H]2-0.51 and E(B—V) from 0.14 to 0.21 (from isochrone fitting).," who quote an age of $\pm$ 1 Gyr and $(m-M)_{0}$ =13.2, with $-$ 0.51 and $E(B-V)$ from 0.14 to 0.21 (from isochrone fitting)."343 Combining the distance modulus with the Galactic coordinates. an A« of about 10 Kpe is obtained.," Combining the distance modulus with the Galactic coordinates, an $R_{\rm gc}$ of about 10 kpc is obtained."344 Previous studies of the metallicity of the cluster are those by Twarog et al. (1995)), Previous studies of the metallicity of the cluster are those by Twarog et al. \cite{twarog95}) )345 who derived [Fe/H]=—0.39 from UBV photometry of TO stars. and by Friel et al. (20025) ," who derived $-$ 0.39 from $UBV$ photometry of TO stars, and by Friel et al. \cite{friel02}) )"346who obtained ]=-0.4720.09 from low resolutior spectroscopy of four cluster giants., who obtained $-0.47\pm0.09$ from low resolution spectroscopy of four cluster giants.347 The only high resolutior spectroscopic investigation carried out so far for Mel 66 is by Gratton Contarint (1994). who studied two giants and obtainec [Fe/H]=—0.38+ 0.15.," The only high resolution spectroscopic investigation carried out so far for Mel 66 is by Gratton Contarini \cite{gratton94}) ), who studied two giants and obtained $=-0.38\pm0.15$ ."348 The properties of the target clusters are summarized i Table 1:: references for ages. distance moduli and reddening are also given; the Galactocentric radii were taken from the collections by Friel et al. (," The properties of the target clusters are summarized in Table \ref{cluster_par}: : references for ages, distance moduli and reddening are also given; the Galactocentric radii were taken from the collections by Friel et al. ("3492002) and Friel (2006): slightly different values from other literature sources were mentioned above.,2002) and Friel (2006); slightly different values from other literature sources were mentioned above.350 The [Fe/H] values cover the whole range found in literaure (see above)., The [Fe/H] values cover the whole range found in literaure (see above).351instability at work. which Leads ultimately to the complete clestruction of the disc.,"instability at work, which leads ultimately to the complete destruction of the disc."352 In Figs., In Figs.353 10. and 11. we plot the time-evolution of the black hole mass and the disc mass for the same set of models displayed in Figs., \ref{fig:Instability3} and \ref{fig:Instability4} we plot the time-evolution of the black hole mass and the disc mass for the same set of models displayed in Figs.354 SN. and 9.. respectively.," \ref{fig:Instability1} and \ref{fig:Instability2}, respectively."355 Ehe sudden loss of the mass of the disc at late times is reflected. on the corresponding rapid increase of the mass of the black hole., The sudden loss of the mass of the disc at late times is reflected on the corresponding rapid increase of the mass of the black hole.356" As an example. for Model (2a). at {ο40fa, the black hole has almost cloubled its mass CAMpg4.7 AL.) aud. correspondingly. the mass of the disc has decreased. from 2.5M. to roughly 0.3M.."," As an example, for Model (2a), at $t\sim 40\ t_{\rm orb}$ the black hole has almost doubled its mass $M_{\rm BH}\sim 4.7\ \mathrm{M}_{\odot}$ ) and, correspondingly, the mass of the disc has decreased from $2.5\ \mathrm{M}_{\odot}$ to roughly $0.3\ \mathrm{M}_{\odot}$."357" Note that since models b and ""€ have a much smaller disc mass than models ‘a’ the growth of the black hole mass in Fig.", Note that since models `b' and `c' have a much smaller disc mass than models `a' the growth of the black hole mass in Fig.358 11 is not as clearly visible as in Fie. 10.., \ref{fig:Instability4} is not as clearly visible as in Fig. \ref{fig:Instability3}.359 The morphology changes that. the. unstable svstenm undergoes ave shown in Figure 12. for à representative case (Model 3a)., The morphology changes that the unstable system undergoes are shown in Figure \ref{fig:HydroModel3a} for a representative case (Model 3a).360 Ehe evolution is qualitatively similar for all models., The evolution is qualitatively similar for all models.361 In this figure we show cight snapshots of the time-evolution from /=θ to /=ILSfaa., In this figure we show eight snapshots of the time-evolution from $t=0$ to $t=11.8\ t_{\rm orb}$.362 The variable plotted in the figure is the rest-mass density., The variable plotted in the figure is the rest-mass density.363 The contour levels are linearly spaced. with Ap=0.1pé. where p? is the maximum value of the densitv at. the center of the initial disc.," The contour levels are linearly spaced with $\Delta\rho=0.1\ \rho_\mathrm{c}^{0}$, where $\rho_\mathrm{c}^{0}$ is the maximum value of the density at the center of the initial disc."364 In Fie., In Fig.365 12. one can clearly follow the transition from a quasistationary aceretion regime (panels (1). to (5)) to the rapid. development of the runaway instability (panels (6) to (8))., \ref{fig:HydroModel3a} one can clearly follow the transition from a quasi-stationary accretion regime (panels (1) to (5)) to the rapid development of the runaway instability (panels (6) to (8)).366 At /=1ISOfa... the disc has almost entirely clisappearecl inside the black hole whose size has noticeably grown.," At $t=11.80\ t_\mathrm{orb}$, the disc has almost entirely disappeared inside the black hole whose size has noticeably grown."367 From the numerical point of view. and as alreacy pointed out in Section 5 when describing stationary models. the flow solution remains considerably smooth even though the evolution is now dynamic.," From the numerical point of view, and as already pointed out in Section \ref{tests} when describing stationary models, the flow solution remains considerably smooth even though the evolution is now dynamic."368 The equatorial plane svinmetry is maintained during the whole evolution. with no sign of numerical asvmumetries as well as no vortices appearing inside the disc., The equatorial plane symmetry is maintained during the whole evolution with no sign of numerical asymmetries as well as no vortices appearing inside the disc.369" Correspondingly. Figure 13. shows the velocity field for model (3a) at f=10.700,45. associated with snapshot (7) in Fie. 12))."," Correspondingly, Figure \ref{fig:HydroModel3a_velocity} shows the velocity field for model (3a) at $t=10.70 t_{\rm orb}$, associated with snapshot (7) in Fig. \ref{fig:HydroModel3a}) )."370 This figure shows that the disc is falling racially. on to the black hole with no signs of vortices and circulation patterns developing., This figure shows that the disc is falling radially on to the black hole with no signs of vortices and circulation patterns developing.371 An interesting information which our hycrodvnamical simulations provide is the timescale of the instability. Fi., An interesting information which our hydrodynamical simulations provide is the timescale of the instability $t_\mathrm{run}$.372 We estimate this timescale as the time it takes for hall o£ the mass of the disc to fall into the hole., We estimate this timescale as the time it takes for half of the mass of the disc to fall into the hole.373 Phe values of fu obtained for our 8 models are given in the last column of ‘Table 3.., The values of $t_\mathrm{run}$ obtained for our 8 models are given in the last column of Table \ref{tab:ModelParameters}.374 Such values span the interval 3.8. 1ἐν which corresponds to very small clurations (6 300ms).," Such values span the interval $\sim 3.8$ $140 t_\mathrm{orb}$, which corresponds to very small durations $\sim 6$ $300\ \mathrm{ms}$ )."375" To check he quality of our definition of fu we have performed the ollowing test: for the four models ofseries ""à we have carried out additional simulations in which the mass of the black iole starts to increase only once the stationary regime has on reached at time /2fy.", To check the quality of our definition of $t_\mathrm{run}$ we have performed the following test: for the four models of series `a' we have carried out additional simulations in which the mass of the black hole starts to increase only once the stationary regime has been reached at time $t \simeq t_{0}$.376 The corresponding evolution of he mass Hux in case (da). for which fo21.9£44. is plotted in Fig. 14...," The corresponding evolution of the mass flux in case (4a), for which $t_{0}\simeq 7.9\ t_\mathrm{orb}$, is plotted in Fig. \ref{fig:TestTimeScale}."377 To compare more easily the timescale associated with the runaway instability in the two cases. we have slotted in Fig.," To compare more easily the timescale associated with the runaway instability in the two cases, we have plotted in Fig."378 15 the evolution of the mass [lux as a function offfy., \ref{fig:TestTimeScaleSerie1} the evolution of the mass flux as a function of $t-t_{0}$.379 Again we see that in all cases the runaway instability appears immediately resulting in the rapid disappearance of he disc., Again we see that in all cases the runaway instability appears immediately resulting in the rapid disappearance of the disc.380 However the precise comparison between our usual series of runs (where fy= 0) and the mocified ones Ieads to he conclusion that the timescale of the runaway instability is a bit overestimated in the first case., However the precise comparison between our usual series of runs (where $t_{0}=0$ ) and the modified ones leads to the conclusion that the timescale of the runaway instability is a bit overestimated in the first case.381 “Phe new values of he timescale for series ‘a’ are reported in the last column of Table 3.., The new values of the timescale for series `a' are reported in the last column of Table \ref{tab:ModelParameters}.382 In Fig., In Fig.383 16. we plot the timescale fu as a function of the mass flux in the stationary regime mu., \ref{fig:TimeScale} we plot the timescale $t_\mathrm{run}$ as a function of the mass flux in the stationary regime $\dot{m}_\mathrm{stat}$.384 To be able to cerive an empirical law for the timescale of the instability. one should. consider a much larger sample of models., To be able to derive an empirical law for the timescale of the instability one should consider a much larger sample of models.385" Nevertheless. two clear tendencies can already be extracted from this figure: (1) the timescale depends weakly on the clise-to-hole mass ratio: (ii) the runaway instability occurs faster when the initial mass flux (stationary value) is larger. following approximatively faixmP""."," Nevertheless, two clear tendencies can already be extracted from this figure: (i) the timescale depends weakly on the disc-to-hole mass ratio; (ii) the runaway instability occurs faster when the initial mass flux (stationary value) is larger, following approximatively $t_\mathrm{run} \propto \dot{m}^{-\alpha}$."386 With the value of a=0.9 obtained. for series ca (where we have used the result of the moclified runs to get a more accurate estimate of £44). we can infer that. for all cases. the disc is cdestroved in a duration never exceeding 1s for a large range of accretion mass fluxes. Heiz107AL.fs.," With the value of $\alpha=0.9$ obtained for series `a' (where we have used the result of the modified runs to get a more accurate estimate of $t_\mathrm{run}$ ), we can infer that, for all cases, the disc is destroyed in a duration never exceeding $1\ \mathrm{s}$ for a large range of accretion mass fluxes, $\dot{m}_\mathrm{stat} \ga 10^{-3}\ \mathrm{M_{\odot}/s}$."387 We have presented. results from a numerical study. of the runaway instability of thick discs around. black holes., We have presented results from a numerical study of the runaway instability of thick discs around black holes.388 In this study we have carried out a comprehensive set of time- simulations aimed at exploring the appearance of the instability., In this study we have carried out a comprehensive set of time-dependent simulations aimed at exploring the appearance of the instability.389 In. order to do so we have used a Lully relativistic. axisvmametric hverodyvnamics code.," In order to do so we have used a fully relativistic, axisymmetric hydrodynamics code."390 The general, The general391 (Gohrenetal.1981). (Filippeuko&IIo2003:Shieldsetal.2008) (Bartheta.200E:Reinesal.2011)... (c.g..Reinestal.," \citep{Gehren84}, \citep{Filippenko03,Shields08} \citep{Barth04,Reines11}. \citep[e.g.,][]{Reines11}."3922011). ioc6 (aia&Loeb2001)., $z\simeq 6$ \citep{Haiman01}.393". s(t~15230) (\""'olonterietal.2003)."," $z \sim39415-30$ \citep{Volonteri03}."395. properties., properties.396" It is conunoulv speculated that Population III stars have a top-heavy IMF. with masses raugnmg from 100 - 1000 (οι,Couchiman&Rees1986:Abel 2001)."," It is commonly speculated that Population III stars have a top-heavy IMF, with masses ranging from 100 - 1000 \citep[e.g.][]{Couchman86,Abel02,Bromm04}."397. Αν zevo-metallicity star wit ja nass greater than ~260 will leave a ~LOOX black hole behind (Bondctal.1981:οσον&Woosley 2002).," Any zero-metallicity star with a mass greater than $\sim 260$ will leave a $\sim 100$ black hole behind \citep{Bond84,Heger02}."398. Another theory involves he direct collapse of very ueal-poor. low-augular momenta gas via ανασα] instabilities (Oh&Taiman2002:LochRasio1991:nanetal.2006:Loato&Natarajan 2006).," Another theory involves the direct collapse of very metal-poor, low-angular momentum gas via dynamical instabilities \citep{Oh02,Loeb94,Eisenstein95,Koushiappas04,Begelman06,Lodato06}."399. If[: enough eas is fueled ito he center of a local overdeusity. it nay collapse to form a black hole with mass 10!10 (Beegchnanctal.2006:Lodato&Natarajan2006:Begchnanetal.2008 ).," If enough gas is funneled into the center of a local overdensity, it may collapse to form a black hole with mass $10^4 - 10^6$ \citep{Begelman06,Lodato06,Begelman08}."400. This process may happen ater than Population III sav formation. because halos uust be larger to host sich iuassive inflow.," This process may happen later than Population III star formation, because halos must be larger to host such massive inflow."401 Effiieut eas collaose ds amore likely to occur dmi iassive halos with virial teniperaures Ju> LOK under ietal-ree conditions where the onuatiou of Πο is inhibited woa UV backeround (Brom&Loch2003). ancl cooling is dominated by atomic bydrogen.," Efficient gas collapse is more likely to occur in massive halos with virial temperatures $T_{\rm vir} > 10^4$ K under metal-free conditions where the formation of $_2$ is inhibited by a UV background \citep{BrommLoeb2003}, and cooling is dominated by atomic hydrogen."402" Iu such halos. racnientation Is stpsressed., cooling proceeds gradually. and the gaseous courponen can cool and participate iu AIDII formation before it is med imto stars."," In such halos, fragmentation is suppressed, cooling proceeds gradually, and the gaseous component can cool and participate in MBH formation before it is turned into stars."403 These halos nav need ο exist nui reelous of ultracritical UV radiation in order to forma MDIIS by «irect collapse (Dijkstraetal. 2010)..since the average estimated UV backeround may not be sufficient toprevent some Population III stars from ornüug iu halos of this size (Johnsonetal. 2008)..," These halos may need to exist in regions of ultracritical UV radiation in order to form MBHs by direct collapse \citep{Dijkstra08,Shang10}, ,since the average estimated UV background may not be sufficient toprevent some Population III stars from forming in halos of this size \citep{Johnson08}. ."404SGT can coexist wilh a nonlinear process active al high E>E.,SGT can coexist with a nonlinear process active at high $E \ge E_{c}$.405 In this case. the PlogE) distribution is modified in two characteristic wavs [rom (he lognormal predicted for pure SGT: G) i£ a decay. process removes energy [rom (he primary (parent) waves. then the P(logE) distribution is reduced near and above £. with known form (Robinsonetal.Robinson1995:Cairus&Grubits2001:CairnsMenietti 2001): Gi) if à nonlinear sell-Iocusing process like wave collapse or modulational instability (Robinson1997). is active. increasing the number of intense wavepackets. then the field distribution is enhanced at hieh E>BÀ. intoa power-law tail with index given in Table L..," In this case, the $P(\log E)$ distribution is modified in two characteristic ways from the lognormal predicted for pure SGT: (i) if a decay process removes energy from the primary (parent) waves, then the $P(\log E)$ distribution is reduced near and above $E_{c}$ with known form \citep{retal1993,r1995,cg2001,cm2001}; (ii) if a nonlinear self-focusing process like wave collapse or modulational instability \citep{r1997} is active, increasing the number of intense wavepackets, then the field distribution is enhanced at high $E \ge E_{c}$ intoa power-law tail with index given in Table \ref{table_theory}."406 Thermal waves driven by an instability and subject to limited stochastic growth effects (but which have not evolved into a pure SGT state) have approximately power-law statistics al fields much greater than the average thermal level Ej (Robinson1995:Cairnsetal.200011: ihe index a in (2)) is positive. depends on the difference between (he average erowth aud damping rates divided by a stochastic parameter. and can (ake a wide range of possible values.," Thermal waves driven by an instability and subject to limited stochastic growth effects (but which have not evolved into a pure SGT state) have approximately power-law statistics at fields much greater than the average thermal level $E_{T}$ \citep{r1995,cetal2000}: the index $\alpha$ in \ref{power-law}) ) is positive, depends on the difference between the average growth and damping rates divided by a stochastic parameter, and can take a wide range of possible values."407 Buportantly. though. the majoritw of the fields will be within a few decades of Ej.," Importantly, though, the majority of the fields will be within a few decades of $E_{T}$."408 Gaussian intensi(v statistics result [rom superposition of multiple random signals. as expected for measurement noise. skv background. and multiple unresolved subsources: In additüon. closely Gaussian intensity statistics can result. [rom scattering of radiation by density inhomogeneities between thie source and observer under some circumstances (Iatcliffe1956:Salpeter1967:BRückett. 1977).," Gaussian intensity statistics result from superposition of multiple random signals, as expected for measurement noise, sky background, and multiple unresolved subsources: In addition, closely Gaussian intensity statistics can result from scattering of radiation by density inhomogeneities between the source and observer under some circumstances \citep{ratcliffe1956,salpeter1967,rickett1977}."409. Refractive locusine can also lead to “lensing” events associated with causties and other singularities (ALA. Walker. personal communication. 2003).," Refractive focusing can also lead to “lensing” events associated with caustics and other singularities (M.A. Walker, personal communication, 2003)."410 In particular. Walker pointed out that the maenilication depends on the geometric properties of caustics: he used established results from gravitational lensing theory [Equation (11.64a) of Schneider et al. (," In particular, Walker pointed out that the magnification depends on the geometric properties of caustics: he used established results from gravitational lensing theory [Equation (11.64a) of Schneider et al. ("4111992)] to show that the distribution of flix magnification factors ji=FfFy lor the lowest order critical curve (a fold) is power-law at large ji. with where Fy is the original flux.,"1992)] to show that the distribution of flux magnification factors $\mu = F/F_{0}$ for the lowest order critical curve (a fold) is power-law at large $\mu$, with where $F_{0}$ is the original flux."412 Convolving this distribution with an initial distribution of Πακος leads to the result (1)) with a=3 al large F'., Convolving this distribution with an initial distribution of fluxes leads to the result \ref{flux_equation}) ) with $\alpha = 3$ at large $F$.413 Independently. the Releree pointed out that a direct analysis of lensing effects in scattering theory [Equation (28a) of (1967)]] vields the form (1)) with a=3.," Independently, the Referee pointed out that a direct analysis of lensing effects in scattering theory [Equation (28a) of \citet{salpeter1967}] ] yields the form \ref{flux_equation}) ) with $\alpha = 3$."414" Predictions exist for the statistics of ""elementary burst” systems (Robinsonetal.1996) and uniform secular growth (Cairns&Robinson 1999)..", Predictions exist for the statistics of “elementary burst” systems \citep{rsw1996} and uniform secular growth \citep{cr1999}. .415 However. these are considered very unlikelv to be relevant and are not detailed here.," However, these are considered very unlikely to be relevant and are not detailed here."416Finally. we consider the constraints placed on the mass-to-light ratio of he mass concentration.,"Finally, we consider the constraints placed on the mass-to-light ratio of the mass concentration."417 Phe measured. flux withins the aperture is. not significantlv.⋠⋠⋅ in. excess of that [οιuic in a sample of control apertures placed randomly on chios 2.3. ancl 4 (avoiding Abell 1942 on chip 1)," The measured flux within the aperture is not significantly in excess of that found in a sample of control apertures placed randomly on chips 2,3, and 4 (avoiding Abell 1942 on chip 1)."418 We use the lo fluctuations in the flux in the control apertures to place an upper limit on the excess luminosity. whic.1 will be dependent on the cosmology and the unknown recshift of the dark clump.," We use the $1\sigma$ fluctuations in the flux in the control apertures to place an upper limit on the excess luminosity, which will be dependent on the cosmology and the unknown redshift of the dark clump."419 Phe total apparent tus is transformc into a total luminosity. using the Alenanteau k-corrections for earlv-tvpe galaxies.," The total apparent flux is transformed into a total luminosity, using the Menanteau $k$ -corrections for early-type galaxies."420 This measurement of Lg is then transformed. into he bolometric luminosity Lua using bolometric corrections for απ old stellar population from the same models., This measurement of $L_H$ is then transformed into the bolometric luminosity $L_{bol}$ using bolometric corrections for an old stellar population from the same models.421 This curve. which varies with the unknown redshift of the clump. is srown in Fig.," This curve, which varies with the unknown redshift of the clump, is shown in Fig."422 7 and. allows comparison with the curves derived in a similar fashion from the {ρα cata., \ref{fig-ml} and allows comparison with the curves derived in a similar fashion from the $I$ -band data.423 As these are upper limits on the luminosity of the clump. the f£-bai data provides tighter constraints on the luminosity at high redshift.," As these are upper limits on the luminosity of the clump, the $H$ -band data provides tighter constraints on the luminosity at high redshift."424" The weak lensing measurements of Erben provide an estimate of the lensing mass within the aperture. assuming a redshift distribution x27οχρί(z/z0)7) for the source galaxies. with £2,?=0.8 ancl ές=1.0 (where 2,5221 ορ]. "," The weak lensing measurements of Erben provide an estimate of the lensing mass within the aperture, assuming a redshift distribution $\propto z^2 \exp(-(z/z_0)^{1.5})$ for the source galaxies, with $\left< z_s \right>=0.8$ and $\left< z_s \right>=1.0$ (where $\left<425z_s \right>\simeq 1.5z_0$ )."426Combining these lensing mass estimates with the upper limits on the bolometric luminosity due to the cluster galaxies. we obtain the lower limits on the MLg; shown in the bottom panel of Fig. 7..," Combining these lensing mass estimates with the upper limits on the bolometric luminosity due to the cluster galaxies, we obtain the lower limits on the $M/L_{Bol}$ shown in the bottom panel of Fig. \ref{fig-ml}."427 From our infrared study we have discounted the possibility that the mass concentration detected by Erben is the result of a hitherto undetected massive high redshift cluster of galaxies., From our infrared study we have discounted the possibility that the mass concentration detected by Erben is the result of a hitherto undetected massive high redshift cluster of galaxies.428 “Phe resulting velocity dispersion. σε. implied bv the measured tangential shear is in excess of 2000 kin + for a singular isothermal sphere model with teu.)= 1.0.," The resulting velocity dispersion, $\sigma_v$, implied by the measured tangential shear is in excess of 2000 km $^{-1}$ for a singular isothermal sphere model with $\left<z_{\rm lens}\right>=1.0$ ."429 This is considerably higher than any currently measured rich cluster. ancl we have shown in Fig.," This is considerably higher than any currently measured rich cluster, and we have shown in Fig."430 1. tha we woulcl expect to make a significant detection of such a cluster to the magnitude imit of our infrared study., \ref{fig-contrast} that we would expect to make a significant detection of such a cluster to the magnitude limit of our infrared study.431 Both searches for an excess in he 4ff colour-magnitude diagram and a spatial overdensity of sources in the deep infrared data reveal no such features., Both searches for an excess in the $I-H$ colour-magnitude diagram and a spatial overdensity of sources in the deep infrared data reveal no such features.432 Furthermore. a massive cluster at higher redshift (bevond the magnitude limit of our study) is ruled out by the lensing constraint. provided by the finite redshift of the background. sources (for which a reasonable limit is," Furthermore, a massive cluster at higher redshift (beyond the magnitude limit of our study) is ruled out by the lensing constraint provided by the finite redshift of the background sources (for which a reasonable limit is"433Second. we caleulated the same intervals for all isochrones on the grid. and selected oulv those with values bius within certain tolerances of the cluster measurements.,"Second, we calculated the same intervals for all isochrones on the grid, and selected only those with values lying within certain tolerances of the cluster measurements."434 In this paper. we adopted 0.2 mae for Amps and £0.02 for Aesμις W," In this paper, we adopted $\pm 0.2$ mag for $\Delta m_{\rm F814W}$ and $\pm 0.02$ for $\Delta c_{m_{\rm F606W}-m_{\rm435F814W}}$."436e fit the selected isochrones to the CMD by eve., We fit the selected isochrones to the CMD by eye.437 At the same time. we calculated the offsets in magnitude aud color required to align the turn-off of the isochrone with that of the CAD. and the offsets iu magnitude required to align the tieht chuup of red IB of the isochrone with that of the CMD. and the offsets in color required to align the RGB fiducial point of the isochrone with that of the CMD.," At the same time, we calculated the offsets in magnitude and color required to align the turn-off of the isochrone with that of the CMD, and the offsets in magnitude required to align the tight clump of red HB of the isochrone with that of the CMD, and the offsets in color required to align the RGB fiducial point of the isochrone with that of the CMD."438 We then averaged the offsets in magnitude aud iu color and applied to overplot the isochroue on the CMD. and identified the best fitting isochrone by eve.," We then averaged the offsets in magnitude and in color and applied to overplot the isochrone on the CMD, and identified the best fitting isochrone by eye."439 The resulting offsets ΗΕ«τν alld OCcoepera Provide estimates for the distance modulus to D379 (Gn AL\y) aud the reddening value (E(B V) Ainmgsiyy=GinMy|chesiw AW ονcopus=-tFouow«Εν.," The resulting offsets $\delta m_{\rm F814W}$ and $\delta c_{m_{\rm F606W}-m_{\rm F814W}}$ provide estimates for the distance modulus to B379 $(m-M)_{0}$ ) and the reddening value $E(B-V)$ ): $\delta m_{\rm F814W}=(m-M)_{0}+A_{\rm F814W}$ , and $\delta c_{m_{\rm F606W}-m_{\rm F814W}}=A_{\rm440F606W}-A_{\rm F814W}$."441 The reddening law from Cardellictal.(1989). is eiiployed in this paper., The reddening law from \citet{car89} is employed in this paper.442 The effective waveleneths of the ACS FOOGW aud ΕσὶΑΝ filters are Age=5918 and 8OG0 (Sirianniotal.2005).. so that from Cardellictal. (1989).. AGGWw=~2SE(BV) aud ARS1IAW~Lav£(BWW) (seeDiriibyctal.2007.fordetails)...," The effective wavelengths of the ACS F606W and F814W filters are $\lambda_{\rm eff}=5918$ and 8060 \citep{sirianni05}, so that from \citet{car89}, , $A_{\rm {F606W}}\simeq 2.8\times E(B-V)$ and $A_{\rm {F814W}}\simeq 1.8\times E(B-V)$ \citep[see][for details]{bm07}."443 The reddening value and distance modulus for D379 obtained in this paper are: E(BV)=0.08 aud (05.M)o=ο where the nnucertaiutv is the staudard error of the mean.," The reddening value and distance modulus for B379 obtained in this paper are: $E(B-V)=0.08$ and $(m-M)_{0}=24.44\pm0.10$ , where the uncertainty is the standard error of the mean."444 The hest-fitting Padova isochrone can be seen in Figure 1l: with the metal abundance 0.009 in Z (or 0.325 in [AL/TI]) aud 11.0 Gyr in age., The best-fitting Padova isochrone can be seen in Figure 1: with the metal abundance $0.009$ in $Z$ (or $-0.325$ in $\rm [M/H]$ ) and 11.0 Gyr in age.445 The age of D379 obtained in this paper is 11.0τε1.5 Cox. where the uncertainty is the standard error of the mean.," The age of B379 obtained in this paper is $11.0\pm1.5$ Gyr, where the uncertainty is the standard error of the mean."446 The priuary purpose of this paper is to obtain the age of D379 bv comparing its CMD with isochirones of the Padova stella: evolutionary models. and to check whether the age of B379 obtained idm this paper. is In aerecineut with the determination of Brownctal.," The primary purpose of this paper is to obtain the age of B379 by comparing its CMD with isochrones of the Padova stellar evolutionary models, and to check whether the age of B379 obtained in this paper, is in agreement with the determination of \citet{brown04a}."447(200 From high-resolution stellar spectroscopy it is presentedla}... that GCs in both the halo aud the bulge of our Galaxy are a/Fe eulianced with the typical values [o/Fo]=0.3dE0.1 dex (seeThomasetal.2003.andreferences therein)..," From high-resolution stellar spectroscopy it is presented that GCs in both the halo and the bulge of our Galaxy are $\rm{\alpha/Fe}$ enhanced with the typical values $\rm{[\alpha/Fe]\approx4480.3\pm0.1}$ dex \citep[see][and references therein]{Thomas03}. ."449 For MBL GCs. the estimates of [a/Fe| ratios bv Beasleyetal.(2005) aud Puziaetal.(2005) showed it may on average be 0.1.0.2 dex lower than in the Milky Way (seealsoColuccietal.2009).," For M31 GCs, the estimates of $\rm{[\alpha/Fe]}$ ratios by \citet{b05} and \citet{ppb05} showed it may on average be $\sim4500.1-0.2$ dex lower than in the Milky Way \citep[see also][]{Colucci09}."451. The Padova stellar evolutionary models do uot provide isochrones with a/Fe|>0.0. however. the luninosities of turn-off. SGB. and of the tip of the ROB are nearly unchanged by varving o chhancement except in the iutermecdiate-age reemnue. where a-cuhanced isochroues are sheltly fainter than scaled solar oues (seeCallartetal.2005.andreferences therein)..," The Padova stellar evolutionary models do not provide isochrones with $\rm{[\alpha/Fe]>0.0}$, however, the luminosities of turn-off, SGB, and of the tip of the RGB are nearly unchanged by varying $\alpha$ enhancement except in the intermediate-age regime, where $\alpha$ -enhanced isochrones are slightly fainter than scaled solar ones \citep[see][and references therein]{Gallart05}. ."452 It is generally kuown that the turn-off. SGB and lower RGB are the most age-seusitive features of the CAID. so. the age of D379 obtained based on the isochrones with |n/Fe|20.0 will not change when using the isochrones with [o/Fe]>0.0.," It is generally known that the turn-off, SGB and lower RGB are the most age-sensitive features of the CMD, so, the age of B379 obtained based on the isochrones with $\rm{[\alpha/Fe]=0.0}$ will not change when using the isochrones with $\rm{[\alpha/Fe]>0.0}$."453 The age of D379 (11.0€1.5 Cyr) obtained in this paper is consistent with the determination (10!1 Car) of Brownetal.(200[2)., The age of B379 $11.0\pm1.5$ Gyr) obtained in this paper is consistent with the determination $10_{-1}^{+2.5}$ Gyr) of \citet{brown04a}.454. Brownotal.(200la} determined he age of D379 by comparing the observed CMD with isochrones of VandeuBereetal.(2006)., \citet{brown04a} determined the age of B379 by comparing the observed CMD with isochrones of \citet{vandenberg06}.455. The result of this per confirmed the couclusion of Brownctal.(200la) that B379 is 23 Cyr vouuger than the oldest Calactic GCs., The result of this paper confirmed the conclusion of \citet{brown04a} that B379 is 2–3 Gyr younger than the oldest Galactic GCs.456 The metallicity of D379 obtained in this paper is 20.325., The metallicity of B379 obtained in this paper is $\rm [M/H]=-0.325$.457 Taking iuto account an cuhaucement ofAL/H| the a-capture elements bv |o/Fe|=0.3 (Brownetal.2001a).. aud using the relation between [MTI]. Fe/TI]. aud [a/Fo] frou Salarisetal.(1993).. we derived Fe/H|=0.5L which is iu good agreement with the determination of =0.53 of Hollandetal.&Carnavich(1908) based ou the shape|Fe/T] of the RGB of the deep CAID observed by the HST/WFEPC2.," Taking into account an enhancement of the $\alpha$ -capture elements by $[\alpha/\rm Fe]=0.3$ \citep{brown04a}, and using the relation between [M/H], [Fe/H], and $[\alpha/\rm Fe]$ from \citet{Salaris93}, we derived $\rm [Fe/H]=-0.54$, which is in good agreement with the determination of $\rm458[Fe/H]=-0.53$ of \citet{hfr97} based on the shape of the RGB of the deep CMD observed by the /WFPC2."459 B379 is located iu the AL31 halo. so the extinction is mainly from the foreground Calactic reddening in the direction of M31. which was discussed by may authors (c.g..vandenDergh1969:MeChure&RacineFrogeletal.1980:FusiPecci 2005)... aud. nearly sinular values were deteriuiuced such as £(5/V)=0.08 o» vau deu Berel (1969). 0.11 bv MeClure&Racine(1969) and Hodge(1992).. 0.08 by Erogeletal.(1980).," B379 is located in the M31 halo, so the extinction is mainly from the foreground Galactic reddening in the direction of M31, which was discussed by many authors \citep[e.g.,][]{vand69,McRa69,fpc80,fusi05}, , and nearly similar values were determined such as $E(B-V)=0.08$ by van den Bergh (1969), 0.11 by \citet{McRa69} and \citet{hodge92}, 0.08 by \citet{fpc80}."460". Iu addition. Barmibyetal.(2000)— determined the reddening for each individual cluster using correlations )etwoeen optical aud infrared colors aud metallicity. aud w defining various ""reddeniug-free parameters using heir large database of iiulti-color photometry."," In addition, \citet{bh00} determined the reddening for each individual cluster using correlations between optical and infrared colors and metallicity, and by defining various “reddening-free” parameters using their large database of multi-color photometry."461 Finally. Barbyetal.(2000). cletermined reddeniugs for 31 clusters. 221 of which are reliable (seeBarbyctal.2000.for details)...," Finally, \citet{bh00} determined reddenings for 314 clusters, 221 of which are reliable \citep[see][for462details]{bh00}."463 For D379. Baruibyetal.(2000.alsoP.Daxiibs.priv.comm.) obtained its reddeniug value o be E(BWV)=0.104 0.05.," For B379, \citet[][also P. Barmby, priv. comm.]{bh00} obtained its reddening value to be $E(B-V)=0.10\pm 0.05$ ."464 It is evident that the reddening value of E(BV)=0.08 obtained iu this xeper is consistent with these determinations., It is evident that the reddening value of $E(B-V)=0.08$ obtained in this paper is consistent with these determinations.465 Given the nuportauce of M21 as an anchor for the extragalactic distance scale; may studies have preseuted distance determinations to N31 using different methods. Pritchet&vande," Given the importance of M31 as an anchor for the extragalactic distance scale, many studies have presented distance determinations to M31 using different methods. \citet{pv87},"466nBergh(1957).. IIollaud(1998). aud Vilardelletal.(2006) have giveu a detailed review.," \citet{holland98}467 and \citet{vilardell06} have given a detailed review."468 Although the stellar populations located in different positions iu M21 have different distance moduli. the dispersion cau be neglected since the distance of M31 is large enough.," Although the stellar populations located in different positions in M31 have different distance moduli, the dispersion can be neglected since the distance of M31 is large enough."469" For example. Richetal.(2005) pointed out that. the clusters in MOI dispersed over a 20 kpc radius would have up to 0.06 mag random distance uncertainty,"," For example, \citet{rich05} pointed out that, the clusters in M31 dispersed over a 20 kpc radius would have up to 0.06 mag random distance uncertainty."470 So. the distance modulus to B379 obtained in this paper should be consisteut with the distance of M31 previously determined within 0.06 mag random distance uncertaintv.," So, the distance modulus to B379 obtained in this paper should be consistent with the distance of M31 previously determined within 0.06 mag random distance uncertainty."471 Now we compared our deteriunuation with the most recent and/or important neasurenmients., Now we compared our determination with the most recent and/or important measurements.472 Freedman&Madore(1990). derived the nean distance modulus to AIBL to be (aM)=διτε0.13 based on the Cepheids iu Daade's fields I. IIT. and IV. (Baade&Swope1963.1965).observed using he Canacda-France-Tawait Telescope CCFITT).," \citet{Freedman90} derived the mean distance modulus to M31 to be $(m-M)_{0}=24.44\pm0.13$ based on the Cepheids in Baade's fields I, III, and IV \citep{bs63,bs65} observed using the Canada-France-Hawaii Telescope (CFHT)."473 Uolland determined thedistauce moduli to 11 AILGCs x fittine theoretical isochrones to the observed. RCBs including D379., \citet{holland98} determined thedistance moduli to 14 M31GCs by fitting theoretical isochrones to the observed RGBs including B379.474 The distance modulus to D379 obtained w Tolland(1998) is (ivMy=21.15+ 0.07.Stanek estimated the distance meoculus o M31as (50.M)=2L171x0.035 by. comparing he red clump stars with parallaxes known to better han in the Hipparcos catalog with the red clump stars in three fields in MOL observed with the HST., The distance modulus to B379 obtained by \citet{holland98} is $(m-M)_{0}=24.45\pm0.07$ \citet{sg98} estimated the distance modulus to M31as $(m-M)_{0}=24.471\pm0.035$ by comparing the red clump stars with parallaxes known to better than in the Hipparcos catalog with the red clump stars in three fields in M31 observed with the .475 Adetermination of Freedmanetal.(2001) based ou, Adetermination of \citet{freedman01} based on476theoretical spectrum would not be appropriate.,theoretical spectrum would not be appropriate.477 Yet. such tests would still serve as sanity checks for the population synthesis computations.," 	Yet, such tests would still serve as sanity checks for the population synthesis computations."478 spectrumm at a reasonably close age to the known one. ie... 41.5 — 4.7 Gyr.," m at a reasonably close age to the known one, i.e., 4.5 – 4.7 Gyr."479 M32 also provides à good test. as we now have resolved visible - IR CMDs that suggest an age of z8.5 Gyr (Grillmair et al.," M32 also provides a good test, as we now have resolved visible – IR CMDs that suggest an age of $\approx 8.5$ Gyr (Grillmair et al."480 1996)., 1996).481 We have adopted the visible — IR. spectrum of M32 from Bruzual Magris (19078)., 	We have adopted the visible – IR spectrum of M32 from Bruzual Magris (1997b).482 Figure 18 shows the fits using the solar abundance models of Yt and of Jimenez (new)., 	Figure 18 shows the fits using the solar abundance models of Yi and of Jimenez (new).483 Both the Yi models and the Jimenez models are based on the theoretical Kuruez spectral library. and thus their fits are limited by the known shortcoming of the Kuruez library in matching the spectra of cool stars.," 	Both the Yi models and the Jimenez models are based on the theoretical Kurucz spectral library, and thus their fits are limited by the known shortcoming of the Kurucz library in matching the spectra of cool stars."484 In addition. these single age. single abundance models may not be good approximations to M32.," 	In addition, these single age, single abundance models may not be good approximations to M32."485 We find that the match becomes better. in particular in the U band and in the IR. when chemically composite models are used.," 	We find that the match becomes better, in particular in the $U$ band and in the IR, when chemically composite models are used."486 Bearing all these limitations in mind. we cautiously find that the Yi models match the M32 spectrum somewhat better.," 	Bearing all these limitations in mind, we cautiously find that the Yi models match the M32 spectrum somewhat better."487 In conclusion. we note the followings.," 	In conclusion, we note the followings."488 The difference in the age estimate between Spinrad et al. (, 	The difference in the age estimate between Spinrad et al. (4891997) and this study is caused by the large difference in the model integrated spectrum.,1997) and this study is caused by the large difference in the model integrated spectrum.490 The Jimenez models. their preferred models. appear much bluer than the Yi and the BC models.," 	The Jimenez models, their preferred models, appear much bluer than the Yi and the BC models."491 The new Jimenez models. presumably improved over his earlier version. are redder than his previous models and thus closer to the Yi and the BC models.," 	The new Jimenez models, presumably improved over his earlier version, are redder than his previous models and thus closer to the Yi and the BC models."492 However. they are still much bluer than the Yi and the BC models.," 	However, they are still much bluer than the Yi and the BC models."493 Currently. the Yi models are in reasonable agreement with the BC models and seem to match the spectra of the sun and M32 at their accepted ages better than the new Jimenez models do.," 	Currently, the Yi models are in reasonable agreement with the BC models and seem to match the spectra of the sun and M32 at their accepted ages better than the new Jimenez models do."494 The pioneering studies of Spinrad. and his collaborators (Dunlop et al., 	The pioneering studies of Spinrad and his collaborators (Dunlop et al.495 1996: Spinrad et al., 1996; Spinrad et al.496 1997) have demonstrated the significance of precise age estimates of distant galaxies., 1997) have demonstrated the significance of precise age estimates of distant galaxies.497" Their suggestion that the red galaxy. S5S3WO9], at z=1.552 is at least 3.5 Gyr old was striking because it would result in a rather strict constraint on cosmology."," 	Their suggestion that the red galaxy, 53W091, at $z = 1.552$ is at least 3.5 Gyr old was striking because it would result in a rather strict constraint on cosmology."498" We have carried out a similar exercise, estimating the age of this galaxy. but only via continuum fitting."," 	We have carried out a similar exercise, estimating the age of this galaxy, but only via continuum fitting."499 When we use the same input parameters. our age estimate is approximately 1.4 — 1.8 Gyr. substantially smaller than theirs. but consistent with those of Bruzual Magris (19973) and of Heap et al. (," 	When we use the same input parameters, our age estimate is approximately 1.4 – 1.8 Gyr, substantially smaller than theirs, but consistent with those of Bruzual Magris (1997a) and of Heap et al. ("5001998).,1998).501 The large age estimate of Spinrad et al., 	The large age estimate of Spinrad et al.502 is apparently caused by the use of the early Jimenez models in their analysis. which are significantly bluer than the Yi and the BC models.," is apparently caused by the use of the early Jimenez models in their analysis, which are significantly bluer than the Yi and the BC models."503 The latest Jimenez models are somewhat redder than his earlier models and closer to the Yi models. resulting in the age estimates that are consistent with our estimates.," 	The latest Jimenez models are somewhat redder than his earlier models and closer to the Yi models, resulting in the age estimates that are consistent with our estimates."504 This may indicate à resolution of the age discrepancy on LBDS 53W091., 	This may indicate a resolution of the age discrepancy on LBDS 53W091.505 We have further improved our estimates over previous ones by adopting convective core overshoot (OS) and realistic metallicity mixtures., 	We have further improved our estimates over previous ones by adopting convective core overshoot (OS) and realistic metallicity mixtures.506 The inclusion of OS has little effect on the UV-based age estimates. but it raises the age estimates based on the visible data normalized to the UV by 20 —50.," 	The inclusion of OS has little effect on the UV-based age estimates, but it raises the age estimates based on the visible data normalized to the UV by 20 –."507. Adopting realistic metallicity distributions is also important because different metallicity groups dominate different parts of the integrated spectrum., 	Adopting realistic metallicity distributions is also important because different metallicity groups dominate different parts of the integrated spectrum.508 If we assume that the majority of stars in 553W09] are already as metal-rich as those, 	If we assume that the majority of stars in 53W091 are already as metal-rich as those509 identified as the ταν liigh state (WITS) (sceΠοιήματαIlis2Os.fordiffereutflavorsofBITstates definitious)..," identified as the very high state (VHS) \citep[see][ for different 510flavors of BH states definitions]{rm,bell05,kw08}."511 Constraimiug the nature of the spectral energy. cutoff| of the power law component is esseutial for understandlue the processes occuriug in the immediate vicinity of he accreting DIT«., Constraining the nature of the spectral energy cutoff of the power law component is essential for understanding the processes occurring in the immediate vicinity of the accreting BHs.512 Until receutly. the observational picture on the extended spectral tails in BID X-ray binaries was based areclv on the OSSE results presented by Croveetal. (1998).," Until recently, the observational picture on the extended spectral tails in BH X-ray binaries was based largely on the OSSE results presented by \citet{grove98}."513".. The main conclusion o| the author Sowas flat in the LIIS the power law has 1 vcleay cutoff at ~ 100 τον, while the IISS the uou-thermal part of the spectrum shows no overturn at lieh enerelies."," The main conclusion of the authors was that in the LHS the power law has a clear cutoff at $\sim$ 100 keV, while the HSS the non-thermal part of the spectrum shows no overturn at high energies."514 The oesencee of: he non-thermal electron populalon Was pro)osed. as all explanation., The presence of the non-thermal electron population was proposed as an explanation.515 The latest observational finding:s. however. reveal much more detailed. pictire of the ctoff energy jehavior.," The latest observational findings, however, reveal much more detailed picture of the cutoff energy behavior."516 Recently Mottactal.(2009) pointed ou je specific pattern iu the high οποιον cutoif evolution imine thiο LIIS-to-IISS spectral transiion in CN 339-|., Recently \citet{motta} pointed out the specific pattern in the high energy cutoff evolution during the LHS-to-HSS spectral transition in GX 339-4.517 Naiuely. the authors reported the momXtonic decrease of jo cuto from 120 keV. in the LIIS o 60 keV du the5 IS aud tjen its sharp increase durius the tr‘ausition Τε) ιο IISS.," Namely, the authors reported the monotonic decrease of the cutoff from 120 keV in the LHS to 60 keV in the IS and then its sharp increase during the transition to the HSS."518 They also pointed out the close connection of ie cutoff energv aud the fast variability iu the source iehteurve., They also pointed out the close connection of the cutoff energy and the fast variability in the source lightcurve.519 The results by Mottaetal.(2009) along with ie cutoff energy behavior in NTE J1550-5(iL reporte iere present much more detailed cutoff phenomenology due to the RATE frequent monitoring capabilitv., The results by \citet{motta} along with the cutoff energy behavior in XTE J1550-564 reported here present much more detailed cutoff phenomenology due to the RXTE frequent monitoring capability.520 These, These521o the cluster background contamination.,to the cluster background contamination.522 Other prominent eatures in the Balmer profiles are the broad blue absorption eatures and the infilline due to emission in the red. wing of the cluster background absorption lines., Other prominent features in the Balmer profiles are the broad blue absorption features and the infilling due to emission in the red wing of the cluster background absorption lines.523 Llovaisky (1989) inds strong red-wing emission in the Balmer lines: his data were obtained on the 3.6-m. Canacla-brance-Llawaii clescope in excellent seeing conditions: therefore the cluster ickeround: Contamination in his spectra is much smaller han in our INT data. where the red-wing emission in the Balmer lines is swamped by the cluster light.," Ilovaisky (1989) finds strong red-wing emission in the Balmer lines: his data were obtained on the 3.6-m Canada-France-Hawaii telescope in excellent seeing conditions; therefore the cluster background contamination in his spectra is much smaller than in our INT data, where the red-wing emission in the Balmer lines is swamped by the cluster light."524 We analysed all the Balmer lines in each spectrum using a multi-Craussian routine. based on Aarquardt’s method of minimization. within the spectral analysis package of Vo dH Mash.," We analysed all the Balmer lines in each spectrum using a multi-Gaussian routine, based on Marquardt's method of minimization, within the spectral analysis package of T. R. Marsh."525 We fitted various combinations of Ciaussians to the line profiles to. try to. model simultaneously. the cluster absorption. the rec wing emission infilling and the broad blue absorption wing.," We fitted various combinations of Gaussians to the line profiles to try to model simultaneously the cluster absorption, the red wing emission infilling and the broad blue absorption wing."526 We found that the models with the best 47 consisted. of an absorption component at. the cluster velocity and two additional bluc-shiltec absorption components., We found that the models with the best ${\chi}^2$ consisted of an absorption component at the cluster velocity and two additional blue-shifted absorption components.527 We find absorption Component 2 to have an average blue shift of ~200 km and —440 kms + for absorption Component 3.," We find absorption component 2 to have an average blue shift of 200 km $^{-1}$, and 440 km $^{-1}$ for absorption component 3."528 Our results are similar to those of Iovaisky. who finds (with data of considerably higher signal-to-noise) that the blue wings of the Balmer lines consist of two or more distinct absorption components with velocities ranging from 150 to SOO km + blue-wards of the cluster background. absorption component.," Our results are similar to those of Ilovaisky, who finds (with data of considerably higher signal-to-noise) that the blue wings of the Balmer lines consist of two or more distinct absorption components with velocities ranging from 150 to 800 km $^{-1}$ blue-wards of the cluster background absorption component."529 The radial velocities of the two blue components are. plotted. in Fig.5.., The radial velocities of the two blue components are plotted in \ref{fig:balmer}.530 What is inumeciately apparent is that there is no obvious orbita phase dependeney., What is immediately apparent is that there is no obvious orbital phase dependency.531 The signal-to-noise of our data is too poor to allow us to investigate the behaviour of the widths of the blue absorption components: in our multi-Ciaussian analysis we fixed the EWLIM of the components using mean values from the Llovaisky elata., The signal-to-noise of our data is too poor to allow us to investigate the behaviour of the widths of the blue absorption components: in our multi-Gaussian analysis we fixed the FWHM of the components using mean values from the Ilovaisky data.532 The Balmer raclial velocities present two challenges: it is dillicult to see how the lines can be free of any phase dependence. and it is dillieult to imagine what could give rise to the two distinct bluc-shifted absorption components GE real).," The Balmer radial velocities present two challenges: it is difficult to see how the lines can be free of any phase dependence, and it is difficult to imagine what could give rise to the two distinct blue-shifted absorption components (if real)."533 Lt is possible that the Balmer absorption profiles represent something more complex than our simple multi-Gaussian models. allow for., It is possible that the Balmer absorption profiles represent something more complex than our simple multi-Gaussian models allow for.534 With so many possible [ree parameters ancl the low signal-to-noise of our line profiles. degeneracy cannot be excluded ancl our minimuni-A7. fits niw not represent the intrinsic nature of the line profiles.," With so many possible free parameters and the low signal-to-noise of our line profiles, degeneracy cannot be excluded and our ${\chi}^2$ fits may not represent the intrinsic nature of the line profiles."535 Lf our racial velocity results are real. then we are dealing with some kind of highly structured mass outllow: the velocities appear to indicate expanding shells or rings.," If our radial velocity results are real, then we are dealing with some kind of highly structured mass outflow: the velocities appear to indicate expanding shells or rings."536 The physical mechanism that might. produce such structures from an rav binary is à nivsterv., The physical mechanism that might produce such structures from an X-ray binary is a mystery.537 Intriguingly. Torres et al. (," Intriguingly, Torres et al. ("5382001) have found. in MMT spectra of C211 obtained in 1996. that the Balmer line velocities do show some evidence of phase dependency. in 1996. and that their greatest blue shift is at phase ~0.25.,"2001) have found, in MMT spectra of AC211 obtained in 1996, that the Balmer line velocities do show some evidence of phase dependency in 1996, and that their greatest blue shift is at phase $\sim 0.25$."539 Other studies (van Zyl et al., Other studies (van Zyl et al.540 in preparation) find that the blue-shifted absorption components of the line profiles in LUSTYSTIS observations of AC211 show no orbital variation in velocity at all., in preparation) find that the blue-shifted absorption components of the line profiles in HST/STIS observations of AC211 show no orbital variation in velocity at all.541 Lt appears that AC211 can exhibit dillerent states. in which the Dalmer line profiles sometimes have an orbital-phase dependence and at other times not.," It appears that AC211 can exhibit different states, in which the Balmer line profiles sometimes have an orbital-phase dependence and at other times not."542 P-Cyveni-profilecl Balmer lines that show no correlation with orbital phase have been observed. in other interacting binaries. for example in the cataclysmic variable BZ Cam. (Ringwalel Naylor 1055) but the reasons for this behaviour are unknown.," P-Cygni-profiled Balmer lines that show no correlation with orbital phase have been observed in other interacting binaries, for example in the cataclysmic variable BZ Cam (Ringwald Naylor 1988), but the reasons for this behaviour are unknown."543 The signal-to-noise in the line is poor. especially in the 1987 and 1988 data with their higher resolution.," The signal-to-noise in the line is poor, especially in the 1987 and 1988 data with their higher resolution."544 We have therefore summed the spectra into (roughly) hour-Iong means., We have therefore summed the spectra into (roughly) hour-long means.545 These hour-long mean spectra do not unfortunately have a signal-to-noise high enough to allow multi-Ciaussian fitting of the line: the two discrete absorption components reported. by Hovaisky. are not. resolved. in. our data., These hour-long mean spectra do not unfortunately have a signal-to-noise high enough to allow multi-Gaussian fitting of the line; the two discrete absorption components reported by Ilovaisky are not resolved in our data.546 We have had to determine the radial velocities by fitting only a single Gaussian to the line., We have had to determine the radial velocities by fitting only a single Gaussian to the line.547 This should nonetheless still give us information about the racial velocity behaviour of the line: Llovaisky’s work shows that one absorption component remains nearly constant., This should nonetheless still give us information about the radial velocity behaviour of the line: Ilovaisky's work shows that one absorption component remains nearly constant.548 “Phis, This549"selected, lower-redshift clusters (??)..","selected, lower-redshift clusters \citep{2011MNRAS.tmp...73A,2010ApJ...718..133H}."550 Our uncertainties are too large to test this for 2270.1., Our uncertainties are too large to test this for 270.1.551" Early cluster formation, consistent with a cluster around 2270.1, would be challenging for cluster formation models, which suggest that massive dark matter halos form primarily at z< 1.2 (?).."," Early cluster formation, consistent with a cluster around 270.1, would be challenging for cluster formation models, which suggest that massive dark matter halos form primarily at $z{_<\atop^{\sim}}$ 1.2 \citep{2007MNRAS.374.1303C}."552" The data are too limited to constrain the X-ray spatial distribution, temperature, or abundances of the diffuse component."," The data are too limited to constrain the X-ray spatial distribution, temperature, or abundances of the diffuse component."553" Our analysis is based on 60 kpc close to or at the center of the cluster, where the physical conditions and the relation to the full R5oo luminosity and temperature depend on the presence/absence of a cooling core (?) and can be disturbed by heat input by the quasar."," Our analysis is based on $\sim 60$ kpc close to or at the center of the cluster, where the physical conditions and the relation to the full $_{500}$ luminosity and temperature depend on the presence/absence of a cooling core \citep{2009A&A...498..361P}554 and can be disturbed by heat input by the quasar."555" Deeper XX-ray observations are required to confirm the extended nature and constrain the spatial distribution, spectral properties, and luminosity of the diffuse X-ray emission."," Deeper X-ray observations are required to confirm the extended nature and constrain the spatial distribution, spectral properties, and luminosity of the diffuse X-ray emission."556" Qualitatively, the observations of extended X-ray emission are consistent with the asymmetric depolarization of the radio source observed by ? (see $3.1)."," Qualitatively, the observations of extended X-ray emission are consistent with the asymmetric depolarization of the radio source observed by \citet{1991MNRAS.250..171G} (see 3.1)."557" For a more quantitative comparison we can consider the simple analysis of ?,, which assumes that the density profile of the cluster can be modeled as a B model and that the energy density in the magnetic field of the depolarizing medium scales as the energy density in thermal particles."," For a more quantitative comparison we can consider the simple analysis of \citet{1991MNRAS.250..198G}, which assumes that the density profile of the cluster can be modeled as a $\beta$ model and that the energy density in the magnetic field of the depolarizing medium scales as the energy density in thermal particles."558" The observed ratio of the dispersions in the Faraday depth in the two lobes, rA in the notation of Garrington Conway, is 2."," The observed ratio of the dispersions in the Faraday depth in the two lobes, $r_\Delta$ in the notation of Garrington Conway, is $\sim 2$."559" Since we detect an excess of counts close to the quasar, the core radius a of any B model that represents the observed X-ray emission must be comparable to or less than the size of our extraction region, with outer radius 7.5"" (64 kpc)."," Since we detect an excess of counts close to the quasar, the core radius $a$ of any $\beta$ model that represents the observed X-ray emission must be comparable to or less than the size of our extraction region, with outer radius $7.5''$ (64 kpc)."560" From the radio data, the projected linear size of each lobe is ((~50 kpc)."," From the radio data, the projected linear size of each lobe is $\sim 50$ kpc)."561 The combination of these values and the observed ΤΑ=2 suggest that the gas distribution is rather flat (B< 0.35) and that the lobes are not aligned close to the line of sight.," The combination of these values and the observed $r_\Delta = 2$ suggest that the gas distribution is rather flat $\beta \la5620.35$ ) and that the lobes are not aligned close to the line of sight."563 However the current quality of X-ray data on the gas near 2270.1 is too low to provide useful constraints., However the current quality of X-ray data on the gas near 270.1 is too low to provide useful constraints.564" Our upcoming, deeper oobservations of this source, orChandra observations of a larger sample of objects with observed depolarization, will give us a probe of the run of gas density and magnetic field strength with radius in high-z clusters which will be difficult to obtain in any other way."," Our upcoming, deeper observations of this source, or observations of a larger sample of objects with observed depolarization, will give us a probe of the run of gas density and magnetic field strength with radius in $z$ clusters which will be difficult to obtain in any other way."565" XX-ray observations of the bright, high-redshift (z=1.532) quasar 2270.1 show strong, unabsorbed emission with a slope I'=1.664:0.08, consistent with expectations for radio-loud quasars."," X-ray observations of the bright, high-redshift $z$ =1.532) quasar 270.1 show strong, unabsorbed power-law emission with a slope $\Gamma = 1.66 \pm 0.08$, consistent with expectations for radio-loud quasars."566 Extended X-ray emission associated with the southern radio lobe of 2270.1 most likely originates in one/both components of the double hotspot within that lobe., Extended X-ray emission associated with the southern radio lobe of 270.1 most likely originates in one/both components of the double hotspot within that lobe.567Spitzer upper limits for the hotspot are inconsistent with synchrotron emission from a single power-law population of electrons as the emission mechanism., upper limits for the hotspot are inconsistent with synchrotron emission from a single power-law population of electrons as the emission mechanism.568" The X-ray emission is consistent with SSC for a magnetic field of nnT, about a third of the equipartition field for that region."," The X-ray emission is consistent with SSC for a magnetic field of nT, about a third of the equipartition field for that region."569 No X-ray emission is detected from the northern radio hotspot as expected from our models based on the observed SED., No X-ray emission is detected from the northern radio hotspot as expected from our models based on the observed SED.570 Faint emission is present just south of the lobe but seems unlikely to be associated given, Faint emission is present just south of the lobe but seems unlikely to be associated given571and moves them slightly closer to the mid-plane.,and moves them slightly closer to the mid-plane.572 This is »ecause dust. settling exageerates the elfects of the planets ον transforming the dise [rom Uarecl to Matter structures., This is because dust settling exaggerates the effects of the planets by transforming the disc from flared to flatter structures.573 ‘Thus. dust settling diminishes contribution from the scaled orque around z=rg because of the higher temperatures.," Thus, dust settling diminishes contribution from the scaled torque around $z=r_H$ because of the higher temperatures."574 At the same time. it strengthens the scaled. torque around he mid-plane region by cust settling as well as planetary eravity.," At the same time, it strengthens the scaled torque around the mid-plane region by dust settling as well as planetary gravity."575 As with the well mixed case. the increment of the orque around. the mid-plane more than compensates the decrement around. z=rg. and consequently dust. settling more accelerates planetary migration (able 3)).," As with the well mixed case, the increment of the torque around the mid-plane more than compensates the decrement around $z=r_H$, and consequently dust settling more accelerates planetary migration (Table \ref{table3}) )."576 Compared with the SOT disc model. the ALAISN disc models tend to move the dips away from the mid-plane while the values around. the mid-plane region slightly. increase.," Compared with the S07 disc model, the MMSN disc models tend to move the dips away from the mid-plane while the values around the mid-plane region slightly increase."577 “These trends are the same as the well mixed case as already discussed., These trends are the same as the well mixed case as already discussed.578 Compared with the well mixed case. however. the scaled. torque dE.dz for the case of dust settling is more sensitive to the change in the background. disc structure.," Compared with the well mixed case, however, the scaled torque $d \Gamma^L_{scal} / dz$ for the case of dust settling is more sensitive to the change in the background disc structure."579 This occurs because dust settling produces cooler and denser dise structures., This occurs because dust settling produces cooler and denser disc structures.580 Fig., Fig.581 11. shows the migration time and rate as a function of orbital radius of a planet., \ref{fig11} shows the migration time and rate as a function of orbital radius of a planet.582 As discussed above. dust settling leads to more rapid migration for both disc models.," As discussed above, dust settling leads to more rapid migration for both disc models."583 Table 3 summarises the dillerence of the migration time between the fully mixed and dust settling cases for both disc models at 1 au., Table \ref{table3} summarises the difference of the migration time between the fully mixed and dust settling cases for both disc models at 1 au.584 Interestinglv. the dillerence between the well mixed and dust settling cases is almost identical for both disce mocels.," Interestingly, the difference between the well mixed and dust settling cases is almost identical for both disc models."585 When the results of the two disc models (MMSN. vs 507) are compared. the dilference diminishes (see the right column of Table 2)).," When the results of the two disc models (MMSN vs S07) are compared, the difference diminishes (see the right column of Table \ref{table2}) )."586 This arises [rom that dust settling gives similar temperature structures for the two cise mocdols. as discussed in & 2.4..," This arises from that dust settling gives similar temperature structures for the two disc models, as discussed in $\S$ \ref{disk_ds}."587 1n this subsection. we briellv ciscuss the combined. effects of planets ancl dust. settling on the torque. because their behaviors are the same as those for the well mixed case.," In this subsection, we briefly discuss the combined effects of planets and dust settling on the torque, because their behaviors are the same as those for the well mixed case."588" We confirmed that the deviation from the function (x L/Ad,,) becomes larger with increasing planeary mass. and hence planetary. migration becomes faster."," We confirmed that the deviation from the function $\propto 1/M_p$ ) becomes larger with increasing planetary mass, and hence planetary migration becomes faster."589 This implics that the distortion created by a planet is simila rto cust settling even though the former is local and the alter is global., This implies that the distortion created by a planet is similar to dust settling even though the former is local and the latter is global.590 Both compress the disc density and lower tje temperature in the mic-plane region., Both compress the disc density and lower the temperature in the mid-plane region.591 Also. we checked that the deviation generally. becomes smaller with increasing the distance rom the central star.," Also, we checked that the deviation generally becomes smaller with increasing the distance from the central star."592 This arises from the combined elfects of dust settling and the power-law structure of the surface density., This arises from the combined effects of dust settling and the power-law structure of the surface density.593 “Phe fatter disc shape. (which is a result of dust settling.) reduces the grazing angle. resulting in a small decrement by the dip formation.," The flatter disc shape, (which is a result of dust settling,) reduces the grazing angle, resulting in a small decrement by the dip formation."594 At the same time. the power-law structure diminishes the ellects of planets as discussed at the well mixed case.," At the same time, the power-law structure diminishes the effects of planets as discussed at the well mixed case."595 We have neelected viscous heating of discs in our mocels., We have neglected viscous heating of discs in our models.596 We have assumed instead that stellar irradiation plays the main role in heating (7)..., We have assumed instead that stellar irradiation plays the main role in heating \citep{cg97}.597 As discussed in Paper LL. viscous heatingὃν dominates stellar irradiation only within 0.1 au for," As discussed in Paper I, viscous heating dominates stellar irradiation only within 0.1 au for"598were accumulated to achieve a 10σ detection threshold of approximately 29 mae (Vega) in the F435W. F606VW. and F715-bands. and approximately 28 mag in the F850LP-band.,"were accumulated to achieve a $10\sigma$ detection threshold of approximately 29 mag (Vega) in the $F435W$, $F606W$, and $F775W$ -bands, and approximately 28 mag in the $F850LP$ -band."599 To further enhance the vast array of deep multi-wavelength imaging for this region. deep near infrared images were also obtained in the UDF target field (IRUDE: GO 9803. HR. Thompson. PI) by acquiring individual pointings with the Near Infrared. Camera ancl Multi Object Spectrometer (NICAIOS) camera 3. each 8 orbits in depth in both the FILOW and FAGOW -bands. (o cover a 2: x 3 mosaic near the center of the UDF.," To further enhance the vast array of deep multi-wavelength imaging for this region, deep near infrared images were also obtained in the UDF target field (IRUDF; GO 9803, R. Thompson, PI) by acquiring individual pointings with the Near Infrared Camera and Multi Object Spectrometer (NICMOS) camera 3, each 8 orbits in depth in both the $F110W$ and $F160W$ -bands, to cover a 3 x 3 mosaic near the center of the UDF."600 This observation was divided into two epochs separated by 90 davs to enable searches lor SNe., This observation was divided into two epochs separated by 90 days to enable searches for SNe.601" Another extremely deep optical field. an ACS ""parallel field”. observed. during the imaging of the IRUDE. overlapped with GOODS South observations."," Another extremely deep optical field, an ACS “parallel field”, observed during the imaging of the IRUDF, overlapped with GOODS South observations."602 The UDF. the UDF ACS Parallel (UDFP). and the IRUDE survevs are unique in that thev are the first surveys sensitive to Type Ia supernovae (SNe Ia) to zx2.2.," The UDF, the UDF ACS Parallel (UDFP), and the IRUDF surveys are unique in that they are the first surveys sensitive to Type Ia supernovae (SNe Ia) to $z \le 2.2$."603 With an elasped time of less than 4 Gyr between the first generation of stars (2~100) and SNe Ia al 2>1.4. the observed number of (hese events al 1.4<2«2.2 could provide an important probe of the assembly time required by SNe la progenitors herealter $04)..," With an elasped time of less than 4 Gyr between the first generation of stars $z \sim 100$ ) and SNe Ia at $z > 1.4$, the observed number of these events at $1.4 < z < 2.2$ could provide an important probe of the assembly time required by SNe Ia progenitors \citep[][hereafter S04]{Strolger2004a}."604 The rate in which SN Ia events occur is governed by the rate in which (heir progenitor stars form. ancl (he (time required for (he SN Ia progenitor to develop into a SN la event.," The rate in which SN Ia events occur is governed by the rate in which their progenitor stars form, and the time required for the SN Ia progenitor to develop into a SN Ia event."605 In the framework of cosmic time. these components are (he star formation rate history. ancl (he population assembly (nme. or the delay Gime function 504)..," In the framework of cosmic time, these components are the star formation rate history, and the population assembly time, or the delay time function \citep[][S04]{1999AA...350..349D}."606" Together. the star formation rate historv and the delay time function describe a model for the SN Ia rate history [R4,(0:)]. which can be compared to observations of the SN Ia rate in different redshift regimes."," Together, the star formation rate history and the delay time function describe a model for the SN Ia rate history $\mathcal R_{Ia}(z)$ ], which can be compared to observations of the SN Ia rate in different redshift regimes."607 In principle. the observed discovery rate of SNe al z>1.4 could help to disünguish between viable models of τίς).," In principle, the observed discovery rate of SNe at $z>1.4$ could help to distinguish between viable models of $\mathcal R_{Ia}(z)$."608 In 904. it was shown that observations of the SN Ia rate over a large redshift range (encompassing the most of the last ~10 billion vears of the universe) are consistent with the combination of the star formation rate history [rom rest-frame Ü-band galaxy studies [hereafter SEI: (2). see Giavaliscoetal. (2004b)]]. and a Gaussian distribution ol delay times. wilh a mean ~4 Gyr and a dispersion of ~1 Gyr.," In S04, it was shown that observations of the SN Ia rate over a large redshift range (encompassing the most of the last $\sim10$ billion years of the universe) are consistent with the combination of the star formation rate history from rest-frame $U$ -band galaxy studies [hereafter $_U$ $z$ ), see \citet{Giavalisco:2003bi}] ], and a Gaussian distribution of delay times, with a mean $\sim 4$ Gyr and a dispersion of $\sim 1$ Gyr."609" llowever. recently it has been shown that the rate SNe la per galaxy al 2<0.1 does appear (o increase towards latergalaxy (vpes. e. ο, Lhe SN Ia rate in irregular (dwarl) galaxies is approximately 10 times lareer than in elliptical galaxies of the same total mass 2004).."," However, recently it has been shown that the rate SNe Ia per galaxy at $z < 0.1$ does appear to increase towards latergalaxy types, e. g. the SN Ia rate in irregular (dwarf) galaxies is approximately 10 times larger than in elliptical galaxies of the same total mass \citep{mannucci}."610 This would seem to imply (hat the SN Ia rate more closely traces (he star formation rate history rather than being largely delayed from it., This would seem to imply that the SN Ia rate more closely traces the star formation rate history rather than being largely delayed from it.611 This is difficult to resolve in light of the GOODS supernova data., This is difficult to resolve in light of the GOODS supernova data.612 ILowever. one could Lwpothesize that the rest-frame U-band galaxy. observations (including dust corrections). in reality. provide a rather poor tracer of actual star formation history. or only a lower limit (Chary&Elbaz 2001)... ," However, one could hypothesize that the rest-frame $U$ -band galaxy observations (including dust corrections), in reality, provide a rather poor tracer of actual star formation history, or only a lower limit \citep{2001ApJ...556..562C}. ."613Or.," Or,"614temperatuce of Ἡον103Ix that the X-ray emission [rom the bubble interior. the shocked wind. itself. becomes significant.,"temperature of $T \sim 10^{4} \K$ that the X-ray emission from the bubble interior, the shocked wind itself, becomes significant."615 This is graphically illustrated by Figs., This is graphically illustrated by Figs.616 SN— and 9.., \ref{fig:t_ew} and \ref{fig:t_rosat}.617 Note how after. shell collapse the averaged: temperature contributing to the photons detected by.ROSAT: Tuosvr- increases again.," Note how after shell collapse the averaged temperature contributing to the photons detected by, $T_{\rm ROSAT}$, increases again."618 This is the emission [rom shocked. wind itself. previously of too low a level compared to the X-ray emission from the shell to be noticeable.," This is the emission from shocked wind itself, previously of too low a level compared to the X-ray emission from the shell to be noticeable."619 9, Fig.620 shows the best fit temperatures [rom single emperature fits to a set of simulatedROSAL spectra over time. for both metal abundance fixed. at. solar. and abundances free to fit.," \ref{fig:t_rosat} shows the best fit temperatures from single temperature fits to a set of simulated spectra over time, for both metal abundance fixed at solar and abundances free to fit."621 As generating LO Poisson realisations or cach of ~50 spectra. fitting ancl finally averaging the results would be prohibitively time consuming. we have itted the spectra with no Poisson noise applied.," As generating 10 Poisson realisations for each of $\sim 50$ spectra, fitting and finally averaging the results would be prohibitively time consuming, we have fitted the spectra with no Poisson noise applied."622 Comparison of this with the more rigorous method. used earlier shows hat this generates essentially identical results. although the values of X7 obtained are misleacingly low.," Comparison of this with the more rigorous method used earlier shows that this generates essentially identical results, although the values of $\chi^{2}$ obtained are misleadingly low."623 As can be seen from Fig., As can be seen from Fig.624 9 the best fit temperatures fail o rellect the temperature changes we know are occurring., \ref{fig:t_rosat} the best fit temperatures fail to reflect the temperature changes we know are occurring.625 Intriguinglv. the best fit metallicity does show a systematic rend between /~5000 and 10000vr.," Intriguingly, the best fit metallicity does show a systematic trend between $t \sim 5000$ and $10626000 \yr$."627 Pwo temperature its. able to fit the shape of the spectrum better. do show an initial drop in both temperatures to a minimum at {55000vr. followed by a increase to a constant value after |9000vr (although the hotter component only levels out after {ο15000 vr)," Two temperature fits, able to fit the shape of the spectrum better, do show an initial drop in both temperatures to a minimum at $t \sim 5000 \yr$, followed by a increase to a constant value after $t628\sim 9000 \yr$ (although the hotter component only levels out after $t629\sim 15 000 \yr$ )."630 To further investigate how spectral fitting to the observed A-ray data depends on the the properties of the bubble. we have repeated the detailed analysis described in Section on the higher wind mass loss rate simulation.," To further investigate how spectral fitting to the observed X-ray data depends on the the properties of the bubble, we have repeated the detailed analysis described in Section \ref{sec:res_t15000} on the higher wind mass loss rate simulation."631 Increasing Aly to 10tALve results in a bubble with the same size at f=15000vr. identical. cold. shell properties. but a shocked wind with a density (temperature) a [actor 2 higher (lower).," Increasing $\Mdot_{\rm W}$ to $10^{-4} \Msol \pyr$ results in a bubble with the same size at $t = 15 000 \yr$, identical cold shell properties, but a shocked wind with a density (temperature) a factor $2$ higher (lower)."632 Single temperature models. give best-fit. results very similar to those given in Table 4 for the Low mass loss rate simulation. except the best fit temperatures are significantIv higher in this case: 7=0.41cx0.09keV. (metal abundance fixed) or Z=0.60+0.13keV. (metal abundance. fitted for).," Single temperature models give best-fit results very similar to those given in Table \ref{tab:rz_fits} for the low mass loss rate simulation, except the best fit temperatures are significantly higher in this case: $T = 0.41\pm{0.09} \keV$ (metal abundance fixed) or $T = 0.60\pm{0.13} \keV$ (metal abundance fitted for)."633 ‘This is despite the average temperature within the shocked. wind being half that of the low mass loss rate simulation., This is despite the average temperature within the shocked wind being half that of the low mass loss rate simulation.634 Again absorption columns (and when fitted for. metallicities) deviate svstematicallv from the true values.," Again absorption columns (and when fitted for, metallicities) deviate systematically from the true values."635 The two temperature spectral models give best. fits very similar to the results given in Table 5.. the best fit temperatures being 71=0.26+0.07keV. and 75=147d0.31 keV.," The two temperature spectral models give best fits very similar to the results given in Table \ref{tab:2rz_fits}, the best fit temperatures being $T_{1} = 0.26\pm{0.07} \keV$ and $T_{2} = 1.47\pm{0.31} \keV$ ."636 Phe only real dillerences are a) the relative [fraction of the total emission measure in the hot. component has increased the spectrum appears to be harder). ancl b) there is evidence for two fit minima. one with low Wy and emission measure. the other a high Ny with high emission Measure.," The only real differences are a) the relative fraction of the total emission measure in the hot component has increased the spectrum appears to be harder), and b) there is evidence for two fit minima, one with low $\nH$ and emission measure, the other a high $\nH$ with high emission measure."637 We shall concentrate on the results from the low mass loss rate simulation. mentioning only where the spectral fit results from the high mass loss rate simulation ciller.," We shall concentrate on the results from the low mass loss rate simulation, mentioning only where the spectral fit results from the high mass loss rate simulation differ."638 Single temperature spectral fits eive misleading information on the state of the bubble. not. surprising given. the multicomponent temperature structure and. the detector Jharacteristics. nor are they sensitive to the true spectral anges that occur during the bubbles growth (see Fig. 9)).," Single temperature spectral fits give misleading information on the state of the bubble, not surprising given the multicomponent temperature structure and the detector characteristics, nor are they sensitive to the true spectral changes that occur during the bubble's growth (see Fig. \ref{fig:t_rosat}) )."639 ‘This suggests that the best-Lit values are strongly allected by 16 multicomponent structure and the instrument response., This suggests that the best-fit values are strongly affected by the multicomponent structure and the instrument response.640 The confidence levels from the fits misrepresent the rue uncertainties in the fitted. parameters., The confidence levels from the fits misrepresent the true uncertainties in the fitted parameters.641 In. particular re Πίος temperatures are apparently well constrained. rough the emission comes from a very broad temperature distribution.," In particular the fitted temperatures are apparently well constrained, although the emission comes from a very broad temperature distribution."642 The spectral fits give little clue to the true emperature distribution. as the fits with metallicity free to it have acceptable 47.2," The spectral fits give little clue to the true temperature distribution, as the fits with metallicity free to fit have acceptable $\chi^{2}$."643 X naive hope that in a situation such as this. with a broad. non-peaked temperature distribution. single temperature fits would show poorly constrained it temperatures is not justified by these results.," A naive hope that in a situation such as this, with a broad, non-peaked temperature distribution, single temperature fits would show poorly constrained best-fit temperatures is not justified by these results."644 Absorbing columns and metal abundances are both systematically mis-fitted., Absorbing columns and metal abundances are both systematically mis-fitted.645 Lt is clear. that by fitting a multicomponent spectrum with a single temperature model we introduce a major systematic effect., It is clear that by fitting a multicomponent spectrum with a single temperature model we introduce a major systematic effect.646 Single temperature its with metal abundance fixed at solar (Table +)) are generally unacceptable. with reduced. 47. ~2. perhaps indicating that the model is not à good. representation of he data.," Single temperature fits with metal abundance fixed at solar (Table \ref{tab:rz_fits}) ) are generally unacceptable, with reduced $\chi^{2} \sim 2$, perhaps indicating that the model is not a good representation of the data."647 The hydrogen columns deviate significantly [ron he true absorbing column., The hydrogen columns deviate significantly from the true absorbing column.648 Freeing the metal abundance to [it does. give. σου its to the data. but at. the expense of giving hest-Lit xwameters that bear little relation to the bubble’s true sroperties.," Freeing the metal abundance to fit does give good fits to the data, but at the expense of giving best-fit parameters that bear little relation to the bubble's true properties."649 “Phe best fit metal abundances are less than one wenticth solar. and apparently. strongly constrained.," The best fit metal abundances are less than one twentieth solar, and apparently strongly constrained."650 This is interesting given the current debate over the accuracy of A-rav determined abundances Bauer Bregman 1996)., This is interesting given the current debate over the accuracy of X-ray determined abundances Bauer Bregman 1996).651 The Εκκοσμος average temperature {οςτς derived in Section 3.2.3. is not à good estimator of the fit temperature. as can be seen in Lig. 9.," The flux-weighted average temperature $T_{\it ROSAT}$ derived in Section \ref{sec:expected_results} is not a good estimator of the best-fit temperature, as can be seen in Fig. \ref{fig:t_rosat}."652 Xn average temperature. even sensibly weighted. does not reflect. how spectral fitting works. and will not give acceptable results.," An average temperature, even sensibly weighted, does not reflect how spectral fitting works, and will not give acceptable results."653 These (Table 5)). give. significantly better fits to. the data than the single temperature mocels and. are slightly better than the dillerential emission measure mocels., These (Table \ref{tab:2rz_fits}) ) give significantly better fits to the data than the single temperature models and are slightly better than the differential emission measure models.654 The absorbing columns fit closer to the true value., The absorbing columns fit closer to the true value.655 The fits are insensitive to the metallicity. and would. not. force us to believe they are significantly cilferent from solar. unlike the single temperature models.," The fits are insensitive to the metallicity, and would not force us to believe they are significantly different from solar, unlike the single temperature models."656"broad line seem to havestronger reflection than those in which the broad, skewed line can be seen.","broad line seem to have reflection than those in which the broad, skewed line can be seen."657" Before drawing this conclusion, however, we consider the fact that the narrow-line observations also exhibit on average a higher disc inclination, (i)=62.9?+7.7? with Model SRI, compared to (i)=34.0°+6.4° for the relativistic observations."," Before drawing this conclusion, however, we consider the fact that the narrow-line observations also exhibit on average a higher disc inclination, $\langle i \rangle = 62.9\de \pm 7.7\de$ with Model SRI, compared to $\langle i \rangle = 34.0\de \pm 6.4\de$ for the relativistic observations."658" This in itself provides a partial explanation for the difficulty in detecting the broad lines in NO7, as the higher inclination sources will suffer stronger Doppler effects, further blending the line into the continuum."," This in itself provides a partial explanation for the difficulty in detecting the broad lines in N07, as the higher inclination sources will suffer stronger Doppler effects, further blending the line into the continuum."659" The fact that the higher inclination sources also show larger reflection fractions is somewhat suspicious, as it implies the total strength of the reflection is similar, whereas it should be weaker in the case of the highly inclined sources."," The fact that the higher inclination sources also show larger reflection fractions is somewhat suspicious, as it implies the total strength of the reflection is similar, whereas it should be weaker in the case of the highly inclined sources."660" To compare these more directly, we use a more robust measure of the reflection strength, which we call the reflected flux fraction (RFF)."," To compare these more directly, we use a more robust measure of the reflection strength, which we call the reflected flux fraction (RFF)."661 This is the ratio of the blurred reflected flux to the continuum flux between 2.5 — 10 keV. The mean RFF of the narrow-line observations is 0.17+0.07 with an intrinsic dispersion of 0.114-0.03., This is the ratio of the blurred reflected flux to the continuum flux between 2.5 – 10 keV. The mean RFF of the narrow-line observations is $\pm$ 0.07 with an intrinsic dispersion of $\pm$ 0.03.662" This is only slightly greater than the mean RFF of the 18 NO7 relativistic-line objects, 0.13+0.03 with an intrinsic dispersion of 0.03+0.01."," This is only slightly greater than the mean RFF of the 18 N07 relativistic-line objects, $\pm$ 0.03 with an intrinsic dispersion of $\pm$ 0.01."663" Therefore, after accounting for inclination effects, the strength of the blurred disc reflection in the narrow-line observations is similar to that of the observations with a relativistically broadened iron line."," Therefore, after accounting for inclination effects, the strength of the blurred disc reflection in the narrow-line observations is similar to that of the observations with a relativistically broadened iron line."664" We have reanalysed the sspectra of 11 Seyfert galaxies classified by NO7 as not requiring a broad, accretion disc iron line in a simple analysis."," We have reanalysed the spectra of 11 Seyfert galaxies classified by N07 as not requiring a broad, accretion disc iron line in a simple analysis."665" By fitting models which allow the disc to be ionised, or to have stronger relativistic effects in a Kerr geometry, we have found that 8 of these 11 observations, in fact, do show evidence for a significant blurred reflection component."," By fitting models which allow the disc to be ionised, or to have stronger relativistic effects in a Kerr geometry, we have found that 8 of these 11 observations, in fact, do show evidence for a significant blurred reflection component."666 The mystery regarding the lack of apparent relativistic reflection in these sources in the N07 analysis can therefore be explained., The mystery regarding the lack of apparent relativistic reflection in these sources in the N07 analysis can therefore be explained.667" The reason is that the most obvious feature in the reflection spectrum — the broad iron Ka line — is rendered indistinguishable from the underlying continuum, by a combination of blending and Comptonisation in an ionised disc, strong relativistic effects and, in some cases, a high disc inclination."," The reason is that the most obvious feature in the reflection spectrum — the broad iron ${\alpha}$ line --- is rendered indistinguishable from the underlying continuum, by a combination of blending and Comptonisation in an ionised disc, strong relativistic effects and, in some cases, a high disc inclination."668" Although the models and parameters are rather degenerate and somewhat difficult to disentangle, the most important effect seems to be that of strong relativistic blurring."," Although the models and parameters are rather degenerate and somewhat difficult to disentangle, the most important effect seems to be that of strong relativistic blurring."669" A mildly ionised reflection component can be detected in MCG+8-11-11, Mrk 6, Mrk 110, NGC 7213 and NGC 7469(1) without the need for a steep emissivity power-law or a spinning black hole."," A mildly ionised reflection component can be detected in MCG+8-11-11, Mrk 6, Mrk 110, NGC 7213 and NGC 7469(1) without the need for a steep emissivity power-law or a spinning black hole."670" On the other hand, seven objects can be well-fit with a cold disc as long as the emission can be concentrated in the very innermost regions where the Kerr metric holds sway."," On the other hand, seven objects can be well-fit with a cold disc as long as the emission can be concentrated in the very innermost regions where the Kerr metric holds sway."671" In reality, it seems that both effects are likely at play in many cases."," In reality, it seems that both effects are likely at play in many cases."672" For example, the strongest detection of blurred reflection is made in Mrk 110 with a combination of both ionisation and strong gravity."," For example, the strongest detection of blurred reflection is made in Mrk 110 with a combination of both ionisation and strong gravity."673" The blurred disc reflection in the narrow-line observations is as strong, if not slightly stronger, as it is in the NO7 observations that are well-fit with a relativistically broadened iron line."," The blurred disc reflection in the narrow-line observations is as strong, if not slightly stronger, as it is in the N07 observations that are well-fit with a relativistically broadened iron line."674" Furthermore, a steep emissivity profile is found in at least five of the observations that give a significant detection of the"," Furthermore, a steep emissivity profile is found in at least five of the observations that give a significant detection of the"675"galaxy, which means metals are lost more efficiently than the gas (H and He).","galaxy, which means metals are lost more efficiently than the gas (H and He)."676 We define the wind “metal-enhanced” when the abundances of metals it carries out are higher than in the ISM., We define the wind “metal-enhanced” when the abundances of metals it carries out are higher than in the ISM.677" Metal-enhanced winds have been already suggested by several dynamical works (e.g. MacLow&Ferrara1999;Recchietal.2001, 2002))."," Metal-enhanced winds have been already suggested by several dynamical works (e.g. \citealt{Mac99,678Recchi01, Recchi02}) )."679 In our models we simply assume a higher wind efficiency weight w; for heavy elements than H and He., In our models we simply assume a higher wind efficiency weight $w_i$ for heavy elements than H and He.680" In particular,we adopt Wi=1(ixH,He), WH,HeX1."," In particular,we adopt $w_i=1 (i\neq\rm H, He)$, $w_{\rm H,He}<1$."681 The wind models with various amounts of metal enhancements are shown in Fig. 10., The wind models with various amounts of metal enhancements are shown in Fig. \ref{Fig:mwdZmuY1}.682". All the models have same input parameters (e=0.5, 7 bursts and the duration is 0.1 Gyr for each one) except for wy He."," All the models have same input parameters $\epsilon=0.5$, 7 bursts and the duration is 0.1 Gyr for each one) except for $w_{\rm H,He}$ ."683" The model experiencing a highly enriched wind (wy,He= 0.1) loses very little gas."," The model experiencing a highly enriched wind $w_{\rm H,He}=0.1$ ) loses very little gas."684" We show also models with mild enhanced wind (wy,He= 0.5) and normal wind (10ΗΗεUo 1)."," We show also models with mild metal-enhanced wind $w_{\rm H,He}=0.5$ ) and normal wind $w_{\rm H,He}=w_{\rm O}=1$ )."685 The evolutionary tracks for the abundance ratios show a loop if the wind is metal-enriched., The evolutionary tracks for the abundance ratios show a loop if the wind is metal-enriched.686" When the wind starts, oxygen is lost more efficiently than hydrogen, hence the oxygen abundance within the galaxy decreases with time in the interburst phase, and elements such as C and Fe will show increasing abundances relative to oxygen owing to their delayed restoration into the ISM."," When the wind starts, oxygen is lost more efficiently than hydrogen, hence the oxygen abundance within the galaxy decreases with time in the interburst phase, and elements such as C and Fe will show increasing abundances relative to oxygen owing to their delayed restoration into the ISM."687 This trend continues until the new burst occurs., This trend continues until the new burst occurs.688" Because of the newly produced oxygen supplied to the ISM the O abundance increases and, as a consequence, the abundances of other elements relative to oxygen decrease."," Because of the newly produced oxygen supplied to the ISM the O abundance increases and, as a consequence, the abundances of other elements relative to oxygen decrease."689" Therefore, the evolutionary track shows a loop."," Therefore, the evolutionary track shows a loop."690" The lower the wH,He, the more the metals lost, the lower the value that the O abundance reaches."," The lower the $w_{\rm H,He}$, the more the metals lost, the lower the value that the O abundance reaches."691" The metal-enhanced wind has also a dramatic influence on the µ—Z relation, as we shown in the upper right panel of Fig. 10.."," The metal-enhanced wind has also a dramatic influence on the $\mu-Z$ relation, as we shown in the upper right panel of Fig. \ref{Fig:mwdZmuY1}. ."692" The normal windmainly reduces the gas fraction rather than the abundance, whereas the"," The normal windmainly reduces the gas fraction rather than the abundance, whereas the"693p=(fh—1)(—DSopp/(GNT)pTUADdHDIE.,"$\rho = \left[\left(h-1\right) \left(\Gamma - 1\right) / 694 \left(K\Gamma\right) \right]^{\left(1/(\Gamma-1)\right)}$."695 Weqwe suppose Chat the gas is non-relativistic.-- choosing D—5/3 and A—0.01.," We suppose that the gas is non-relativistic, choosing $\Gamma = 5/3$ and $K=0.01$."696 Integrating over the volume of the initial gas distribution. we find a tolal rest-mass of 353.," Integrating over the volume of the initial gas distribution, we find a total rest-mass of 353."697 This is 20% larger than that in simulation ADP. a shift due to our slightly different choice of Aij.," This is $20\%$ larger than that in simulation KDP, a shift due to our slightly different choice of $l_\mathrm{in}$."698 Note that the code units of gas mass are completely arbitrary., Note that the code units of gas mass are completely arbitrary.699 The initial magnetic field lies entirely within the torus and follows contours of constant density., The initial magnetic field lies entirely within the torus and follows contours of constant density.700 The magnitude of the magnetic field is set so that the volume-weightecl integrated magnetic pressure is 100 times less than the volume-weighted integrated gas pressure., The magnitude of the magnetic field is set so that the volume-weighted integrated magnetic pressure is $100$ times less than the volume-weighted integrated gas pressure.701 The atmosphere surrounding the disk is unmagnetized and static., The atmosphere surrounding the disk is unmagnetized and static.702" The aümosphere's density ancl pressure are set to their smallest. allowed. values. which are chosen so that the floor state is in approximate pressure equilibrium: pg,=7xLO""puaur and TX10οι""(P— 0). where py; is the initial maxinnumn value of the rest-mass density in the disk."," The atmosphere's density and pressure are set to their smallest allowed values, which are chosen so that the floor state is in approximate pressure equilibrium: $\rho_\mathrm{floor} = 7\times10^{-9}\rho_\mathrm{max} r^{-3/2}$ and $P_\mathrm{floor} = 7\times10^{-11}\rho_\mathrm{max} r^{-5/2}\left(\Gamma-1\right)$ , where $\rho_\mathrm{max}$ is the initial maximum value of the rest-mass density in the disk."703 A magnetized accretion disk is subject to the magneto-rotational instability (ATRL). which transfers angular momentum outward.," A magnetized accretion disk is subject to the magneto-rotational instability (MRI), which transfers angular momentum outward."704 This transfer taps into the available free enerev of differential rotation. creating the magnetic fields aud poloidal velocity fluctuations that make up the resulting MIID turbulence.," This transfer taps into the available free energy of differential rotation, creating the magnetic fields and poloidal velocity fluctuations that make up the resulting MHD turbulence."705 This turbulence is dissipative: magnetic and kinetic enerev is lost numerically at the eridscale., This turbulence is dissipative; magnetic and kinetic energy is lost numerically at the gridscale.706 Equation (7)). however. ensures that in the nunmerical solution all (hat dissipated energy is converted to heat.," Equation \ref{conservative-eq}) ), however, ensures that in the numerical solution all that dissipated energy is converted to heat."707 If that heat were retained bv the (id. the disk would become ever hotter ancl geometrically thicker.," If that heat were retained by the fluid, the disk would become ever hotter and geometrically thicker."708 Ultimately. the thermal energv. would either be acereted bv the hole or be carried out [rom the disk bv a wind., Ultimately the thermal energy would either be accreted by the hole or be carried out from the disk by a wind.709 By adding a loss term (o the energy equation. we can estimate either the luminosity ol those svstems in which radiation is efficient or the total heat generated in those svstems in which it is not.," By adding a loss term to the energy equation, we can estimate either the luminosity of those systems in which radiation is efficient or the total heat generated in those systems in which it is not."710 We assume that (he radiation described by (his loss term is optically (hin., We assume that the radiation described by this loss term is optically thin.711" It therefore acts as à passive sink in the local energy. conservation equation (2)): where F,, is (he amount of radiated energy-momentunm per unit 4-volume in the coordinate frame.", It therefore acts as a passive sink in the local energy conservation equation \ref{energy-conservation-eq}) ): where $\mathcal{F}_\nu$ is the amount of radiated energy-momentum per unit 4-volume in the coordinate frame.712" To describe the radiation. we make the simplest assumption: that (he οΙΡΙΟ. is isolropic in the fluids frame: where the ""cooling Iuuection” £ is (he rate energy is radiated per unit proper time in the fluid frame."," To describe the radiation, we make the simplest assumption: that the emission is isotropic in the fluid's frame: where the “cooling function” $\lum$ is the rate energy is radiated per unit proper time in the fluid frame."713if oue or more periods are conteniporaneouslv excited and deteriuuiug them precise values (and the shape of corresponding lielit or radial velocity curves) is riaudatory to ascertain the nature of the iuntriusic variability iu V132 Aur.,if one or more periods are contemporaneously excited and determining their precise values (and the shape of corresponding light or radial velocity curves) is mandatory to ascertain the nature of the intrinsic variability in V432 Aur.714 To this atu we plan for the coming new observiug season to obtain night-long cousecutive series of photometric (optical aud possibly infrared too) aud radial velocity observations of V132 Aur., To this aim we plan for the coming new observing season to obtain night-long consecutive series of photometric (optical and possibly infrared too) and radial velocity observations of V432 Aur.715in the vicinity of the cyclotron energy. further out in the atinosphere aud thus closer to that of he photons at the cyclotron energy.,in the vicinity of the cyclotron energy further out in the atmosphere and thus closer to that of the photons at the cyclotron energy.716 These sinall equivalent widths may help explain the lack of larrow absorption features in the observatious of anomalous X-ray pulsars aud soft. gaiuma-ray 'epeaters CJuett et 22001: Patel et 22001)., These small equivalent widths may help explain the lack of narrow absorption features in the observations of anomalous X-ray pulsars and soft gamma-ray repeaters (Juett et 2001; Patel et 2001).717 Deeper observations of these sources. especially with instruments that have higher sensitivity al EZ5 keV range may help reveal some of these eatures aud thus the uature of these intriguing objects.," Deeper observations of these sources, especially with instruments that have higher sensitivity at $E \gtrsim 5$ keV range may help reveal some of these features and thus the nature of these intriguing objects."718 IL thank John Bahcall for his valuable input. aud cliscussious ou tle plivsies of mode evolution across resonances., I thank John Bahcall for his valuable input and discussions on the physics of mode evolution across resonances.719 [also thauk Dimitrios Psaltis for many discussions on tlie treatment of resonances iu radiative trausler problems aud Ramesh Naravau for useful suggestious., I also thank Dimitrios Psaltis for many discussions on the treatment of resonances in radiative transfer problems and Ramesh Narayan for useful suggestions.720 This work was supported in part by a fellowship of the Ixeck Fouudatiou aud au NSF graut PHY-0070023., This work was supported in part by a fellowship of the Keck Foundation and an NSF grant PHY-0070928.721 In a recent paper. Lai Ho (2002) discussed the pliysics of vacuum polarization: resonance aud pointed out the well-knowu effects of adiabatic evolution aud enhanced polarization mode conversion that take place through the resonant deusity (see the references iu 1).," In a recent paper, Lai Ho (2002) discussed the physics of vacuum polarization resonance and pointed out the well-known effects of adiabatic evolution and enhanced polarization mode conversion that take place through the resonant density (see the references in 1)."722 They argued that this effect was not treated in previous studies aud evaluated the conditions uuder which the large Faraday depolarization assumption and heuce the adiabatie evolution of mocles breaks down., They argued that this effect was not treated in previous studies and evaluated the conditions under which the large Faraday depolarization assumption and hence the adiabatic evolution of modes breaks down.723 In a subsequent work. Ho Lai (2002) further claimed. that they included for the first time the ellect of this uew pheuoimenon on the spectra of a magnetized neutron star atinosphiere in radiative equilibrium.," In a subsequent work, Ho Lai (2002) further claimed that they included for the first time the effect of this new phenomenon on the spectra of a magnetized neutron star atmosphere in radiative equilibrium."724 Iu this appendix. we clarity the ellects of vacuum polarization resonance on the norimal-1uode description of photon trausport in magnetized inedia.," In this appendix, we clarify the effects of vacuum polarization resonance on the normal-mode description of photon transport in magnetized media."725 We show that the effects discussecl by Lai Ho (2002) have been taken into account in the previous calculations of photou trausport through a plasma (Bulik Miller 1997: Zane et 22000: Ozzel 2001)., We show that the effects discussed by Lai Ho (2002) have been taken into account in the previous calculations of photon transport through a plasma (Bulik Miller 1997; Zane et 2000; Özzel 2001).726 In particular. as long as a treatinent is emmploved. the pliysies included in the calculatious is the same. indepeudent of the nomenclature with which one describes modes above aud below the resonant density.," In particular, as long as a normal-mode treatment is employed, the physics included in the calculations is the same, independent of the nomenclature with which one describes modes above and below the resonant density."727 Below. we outline some of the mistakes aud inconsistencies in their discussion.," Below, we outline some of the mistakes and inconsistencies in their discussion."728" First. referring to the case where the photou mocles evolve adiabatically as ""mode conversion” is tnisleacling. since this is precisely the case where the normal modesof propagation (ἳ=1.2 or —/+ mocles) remains the same above and below the resonance."," First, referring to the case where the photon modes evolve adiabatically as “mode conversion” is misleading, since this is precisely the case where the normal modesof propagation $i=1,2$ or $-/+$ modes) remains the same above and below the resonance."729 What dillerent is the correspondeuce between the polarization auc propagation eigeustates ou the two sides of tlie resonance., What different is the correspondence between the polarization and propagation eigenstates on the two sides of the resonance.730 Note iu particular that equatious (2.27) aud (2.13) of Ho Lai (2002). which presumably describe the two different definitions of normal modes. are mathematically identical.," Note in particular that equations (2.27) and (2.43) of Ho Lai (2002), which presumably describe the two different definitions of normal modes, are mathematically identical."731 Secoud. adiabatic evolution is not au additioual effect that needs to be in calculations," Second, adiabatic evolution is not an additional effect that needs to be in calculations"732neulron star surface with thermal evolution calculations it is possible to investigate (he physical processes (hat occur in the interiors of these objects (Schaaletal.1999).,neutron star surface with thermal evolution calculations it is possible to investigate the physical processes that occur in the interiors of these objects \citep{sch99}.733. Although many rotation-powered pulsus have now been observed in X-ravs. when one considers (he number of known radio pulsars. (he sample of sources detected αἱ energies is still small.," Although many rotation-powered pulsars have now been observed in X-rays, when one considers the number of known radio pulsars, the sample of sources detected at X-ray energies is still small."734 The prospects of detecting a greater number of these objects is however improving with missions like aandChandra., The prospects of detecting a greater number of these objects is however improving with missions like and.735 It is now possible. even for [aint aud distant objects like the one presented in (his paper. to discriminate whether (he dominant mechanism is thermal or non-thermal which in tum permits the first estimates of temperature and/or power law index.," It is now possible, even for faint and distant objects like the one presented in this paper, to discriminate whether the dominant mechanism is thermal or non-thermal which in turn permits the first estimates of temperature and/or power law index."736 In particular. the detecüon of emission due to cooling is essential to determine where the pulsar Dies on (he thermal/evoluGonary diagram. with the main objective being able to constrain observationallv (hie equation of state of matter at supra-nuclear densities.," In particular, the detection of emission due to cooling is essential to determine where the pulsar lies on the thermal/evolutionary diagram, with the main objective being able to constrain observationally the equation of state of matter at supra-nuclear densities."737 Currently. (here are very [few sources for which this is possible. (herelore anv new source that can be added to the parameter space is important.," Currently, there are very few sources for which this is possible, therefore any new source that can be added to the parameter space is important."738 One such object that should fall into the category of pulsars for which the cooling neulron star is the dominant source of emission isB2334+61., One such object that should fall into the category of pulsars for which the cooling neutron star is the dominant source of emission is.739. This Vela-like pulsar is located at a distance of D=3.1n kpe (Cordes&Lazio2002)., This Vela-like pulsar is located at a distance of $D=3.1^{+0.2}_{-1.0}$ kpc \citep{cor02}.740. XAnalvsis by implies that the pulsar is associated with the SNR G114.2. making the source one of the oldest (spin-down age of ~4x10! vr) that is still linked to à SNR.," Analysis by \citet{kul93} implies that the pulsar is associated with the SNR G114.3, making the source one of the oldest (spin-down age of $\sim 4\times 10^{4}$ yr) that is still linked to a SNR."741 Hence. iis a kev object to study with relationto the thermal/evolutionary parameter space.," Hence, is a key object to study with relationto the thermal/evolutionary parameter space."742" The spin-down rate (P—191x10P ss 1] of the 495 ms pulsar indicates that the magnetic field of iis B~10"" G. wwas originally detected in a short ( ks)ROSAT pointing (Beckeretal.1993.1996) but the low statistics prevented a comprehensive investigation of the spectral properties of the source or a meaninglul temporal analvsis."," The spin-down rate $\dot P = 191\times 10^{-15}$ s $^{-1}$ ) of the 495 ms pulsar indicates that the magnetic field of is $B\sim 10^{13}$ G. was originally detected in a short (8 ks) pointing \citep{bec93,bec96} but the low statistics prevented a comprehensive investigation of the spectral properties of the source or a meaningful temporal analysis."743 Here we report on the first oobservation ofD23344-61., Here we report on the first observation of.744. wwas observed with oon 2004 February 12., was observed with on 2004 February 12.745 For the spectral and liming analysis we used data from the European Photon Imaging Camera (EPIC) instruments: the EPIC MOS detector (Turner and the EPIC-PN detector (Strüderetal.2001)., For the spectral and timing analysis we used data from the European Photon Imaging Camera (EPIC) instruments: the EPIC MOS detector \citep{tur01} and the EPIC-PN detector \citep{str01}.746. Both MOS instruments and the PN were configured in frame mode and we used the thin filter., Both MOS instruments and the PN were configured in mode and we used the thin filter.747 The MOS1 and MOS2 observations, The MOS1 and MOS2 observations748parameters using Balmer line profile fitting may well amount to a few hundred degrees.,parameters using Balmer line profile fitting may well amount to a few hundred degrees.749 Fortunately. because ionized calemum is the dominant species. the abundance of calcium based on lines is not sensitive to temperature.," Fortunately, because ionized calcium is the dominant species, the abundance of calcium based on lines is not sensitive to temperature."750 On the other hand. as we have demonstrated. it does show a mild dependence on surface gravity because a higher electron pressure favours neutral calcium.," On the other hand, as we have demonstrated, it does show a mild dependence on surface gravity because a higher electron pressure favours neutral calcium."751 Adopting conservative error bars for the temperature and surface gravity we calculated an absolute V magnitude using the mass-radius relations of Benvenuto&AI-thaus14.9707. (1999)., Adopting conservative error bars for the temperature and surface gravity we calculated an absolute $V$ magnitude $M_V=14.9^{+0.9}_{-0.6}$ using the mass-radius relations of \citet{ben1999}.752". The distance modulus implies à. photometric distance d=391, pe.", The distance modulus implies a photometric distance $d=33^{+9}_{-14}$ pc.753 The object is relatively old with a cooling age foo=3.5—8.0 Gyr. but with an uncertain mass (0.35—0.91 M..).," The object is relatively old with a cooling age $t_{\rm cool}=3.5-8.0$ Gyr, but with an uncertain mass $0.35-0.91\,M_\odot$ )."754 We determined the Galactic velocity vector UVW (Table 5)) using our distance estimate and radial velocity measurement (Sect., We determined the Galactic velocity vector $UVW$ (Table \ref{tbl-prop}) ) using our distance estimate and radial velocity measurement (Sect.755 3.3). and published proper-motion.," 3.3), and published proper-motion."756 We employed the algorithm of Johnson&Soderblom(1987)., We employed the algorithm of \citet{joh1987}.757. The kinematics imply membership to the old thin disk (Sionet1988) consistent with the upper range of our age estimate., The kinematics imply membership to the old thin disk \citep{sio1988} consistent with the upper range of our age estimate.758 The acquisition of broadband UÜBV and JHK photometry and of a parallax measurement should help determine the atmospheric parameters more precisely., The acquisition of broadband $UBV$ and $JHK$ photometry and of a parallax measurement should help determine the atmospheric parameters more precisely.759 The stellar radius. hence surface gravity measurement would be improved with a parallax measurement.," The stellar radius, hence surface gravity measurement would be improved with a parallax measurement."760 Accordingly. the error on the calcium abundance measurement would be reduced.," Accordingly, the error on the calcium abundance measurement would be reduced."761 Accurate J/HK photometry would also allow us to investigate possible infrared excess and the presence of a debris disc., Accurate $JHK$ photometry would also allow us to investigate possible infrared excess and the presence of a debris disc.762 The DAZ white dwarfs G 77-50 and G 174-74 are part of a survey including the coolest known DAZ white dwarfs (Zuckermanetal..2003)., The DAZ white dwarfs G $-$ 50 and G $-$ 74 are part of a survey including the coolest known DAZ white dwarfs \citep{zuc2003}.763. In cool convective white dwarfs. heavy elements diffuse below the mixed convective layers. and their presence in white dwarf atmospheres is transitory.," In cool convective white dwarfs, heavy elements diffuse below the mixed convective layers, and their presence in white dwarf atmospheres is transitory."764 Koester&Wilken(2006) estimated the diffusior time-scale for various heavy elements., \citet{koe2006} estimated the diffusion time-scale for various heavy elements.765" The accretion rate required to sustain a given mass fraction X in the atmosphere is given by where M, Is the mass of the convection zone and r the diffusion. time-scale at the bottom of the convection zone. where diffusion is allowed to take place."," The accretion rate required to sustain a given mass fraction $X$ in the atmosphere is given by where $M_{\rm cvz}$ is the mass of the convection zone and $\tau$ the diffusion time-scale at the bottom of the convection zone, where diffusion is allowed to take place."766" The ratio X/X, 15 the ratio of the measured mass-fraction to the accreted mass fraction.", The ratio $X/X_{\rm acc}$ is the ratio of the measured mass-fraction to the accreted mass fraction.767" Therefore. the mass accretion rate of any particular element Is Adopting. in theT appropriate temperature range. a value for the slow-varying ratio Mi; to diffusion time scale of logMa,τοι)=-115 in units of Ms yyr7! or =14.3 in units of ess”! (Koester&Wilken.2006).. we estimated the mass accretion rate of ealeium (in ss!) to be where Xe,=[Aca(C]/LÀAg1H)]2x107. and Aca and Ay are the atomic weights."," Therefore, the mass accretion rate of any particular element is Adopting, in the appropriate temperature range, a value for the slow-varying ratio $M_{\rm cvz}$ to diffusion time scale of $\log{(M_{\rm cvz}/\tau_{\rm Ca})}\approx -11.5$ in units of $M_\odot$ $^{-1}$ or $=14.3$ in units of $^{-1}$ \citep{koe2006}, we estimated the mass accretion rate of calcium (in $^{-1}$ ) to be where $X_{\rm Ca} \approx [A_{\rm Ca}\,n({\rm Ca})]/[A_{\rm H}\,n({\rm H})]= 2\times10^{-9}$, and $A_{\rm Ca}$ and $A_{\rm H}$ are the atomic weights."768 Assuming calcium is accreted as part of a solar-composition flow. the total mass accretion rate (including hydrogen) is 5x10? gss! or 8x107Ma yyr!.," Assuming calcium is accreted as part of a solar-composition flow, the total mass accretion rate (including hydrogen) is $5\times10^9$ $^{-1}$ or $8\times10^{-17}\ M_\odot$ $^{-1}$."769 Our measured error on the caletum abundance of 40.3 dex translates into a similar error on the calculated aceretion rate onto the white dwarf surface., Our measured error on the calcium abundance of $\pm0.3$ dex translates into a similar error on the calculated accretion rate onto the white dwarf surface.770 The true error may well be much larger., The true error may well be much larger.771 Koester(2009) considers that the application of the mixing-length theory to the structure of convection zones may underestimate the mass of the mixed layers by orders of magnitude., \citet{koe2009} considers that the application of the mixing-length theory to the structure of convection zones may underestimate the mass of the mixed layers by orders of magnitude.772" The effect of ""under-shooting below the convection zone may affect diffusion time-scale estimates.", The effect of “under”-shooting below the convection zone may affect diffusion time-scale estimates.773 Therefore. the precision claimed in measuring abundance of parent bodies may be over-estimated.," Therefore, the precision claimed in measuring abundance of parent bodies may be over-estimated."774 Farihietal.(2011) propose that the model of Potter&Tout(2010) for the presence of a magnetic field in post-common envelope (CE) binaries could also be applied to CE episodes with planetary rather than stellar secondary components., \citet{far2011} propose that the model of \citet{pot2010} for the presence of a magnetic field in post-common envelope (CE) binaries could also be applied to CE episodes with planetary rather than stellar secondary components.775 Whether the magnetic field i5 acquired during such a process. or whether it ts a fossil field cannot be ascertained for individual objects but rather from population studies (see.e.g..Kawka&Vennes.20044:etal..2007:Wickramasinghe&Ferrario. 2005).," Whether the magnetic field is acquired during such a process, or whether it is a fossil field cannot be ascertained for individual objects but rather from population studies \citep[see, e.g.,][]{kaw2004a,kaw2007,wic2005}."776. In the case of NLTT 10480. which is old (>3.5 Gyr) with relatively short diffusion time-scales. the present-day metallicity is not linked to the CE event that potentially generated the magnetic field. but more likely to a recent accretion event.," In the case of NLTT 10480, which is old $>3.5$ Gyr) with relatively short diffusion time-scales, the present-day metallicity is not linked to the CE event that potentially generated the magnetic field, but more likely to a recent accretion event."777 A low incidence of planetary systems would imply a low incidence of weak magnetic fields (B<| MG)., A low incidence of planetary systems would imply a low incidence of weak magnetic fields $B\la 1$ MG).778 Current data indicate a low incidence of weak magnetic fields. and Kawkaetal.(2007) found that 6 out of 53 local white dwarfs (d<20 pe) observed with sufficient accuracy to unveil fields weaker than | MG were found to harbour such a low field.," Current data indicate a low incidence of weak magnetic fields, and \citet{kaw2007} found that 6 out of 53 local white dwarfs $d \le 20$ pc) observed with sufficient accuracy to unveil fields weaker than 1 MG were found to harbour such a low field."779 Joreover. Kawka&Vennes(2004a) found that low-field white dwarfs lack progenitors. a gap that could be filled with the CE-," Moreover, \citet{kaw2004a} found that low-field white dwarfs lack progenitors, a gap that could be filled with the CE-mechanism."780 Early results from the survey also indicate a low incidence of very large planets in short-period orbits (P.<50 , Early results from the survey also indicate a low incidence of very large planets in short-period orbits $P\la 50$ 781 Woosley1993)). (Calamaetal.1998.. Ίνα two (Fruchteretal.1999)) (Bloometal.19993) ~230 Bloomet(1999)... Calamaetal.(1999).. Reichart(1999) Wasian(1999). (AkerlofVetal.1999)) (Mészáros.Rees&Papathanassiou199L.. Alészáros&Rees1997.. Panaitescu&Alészaros1908.. Siri&Piran1999)). Wasiman&Draine ~2300 ~2(1|:)jauu + ~10 0.OSjaa," \fcitep{woo93}) \fcitep{gea98}, \fcitep{kea98}) \fcitep{fea99}) \fcitep{bea99}) $\sim 30$ \fcitet{bea99}, \fcitet{gea99}, \fcitet{rei99} \fcitet{wad99}782 \fcitep{aea99}) \fcitep{mrp94}, \fcitep{mer97}, \fcitep{pam98}, \fcitep{sap99}) \acite{wad99} $\sim2300$ $\sim 2\,(1+z)\,{\rm \mu m}$ $z$ $\sim10$ $0.4-0.8\,{\rm \mu m}$ \\ref{obs}. \\ref{impl}."783 /=0.6.Oxy=0.3. O4=0.7 D4-1.52 0.5S2833. Wasian&Draiue(1999) ,$h=0.6$$\Omega_{\rm M} = 0.3$ $\Omega_{\Lambda} = 0.7$ $D_{\rm A}=1.5-2$ $0.5\lsim z\lsim3$ \cite{wad99} 784Over the past decade our understanding of the structural and physical properties of dises around young stars has increased from basic theoretical modelling of the spectral energy distributions (SEDs) constrained by observations with no spatial information. to modelling based on not only the SEDs. but also spatially resolved dust observations. like scattered light images and interferometry (???).,"Over the past decade our understanding of the structural and physical properties of discs around young stars has increased from basic theoretical modelling of the spectral energy distributions (SEDs) constrained by observations with no spatial information, to modelling based on not only the SEDs, but also spatially resolved dust observations, like scattered light images and interferometry ."785. Two decades ago. the first. submillimetre interferometer observations resolved the molecular gas emission spatially and this allowed major progress in understanding the dise kinematics. structure and chemistry 222?).. (," Two decades ago, the first submillimetre interferometer observations resolved the molecular gas emission spatially and this allowed major progress in understanding the disc kinematics, structure and chemistry . ("786Subjmillimetre gas and dust emission is the ideal probe of the global dise properties. like size. nass and radial distribution of dise material. because the bulk of the dise mass is located beyond 100 AU from the star. at temperatures of 10-50 K that dominate this part of the spectrum.,"Sub)millimetre gas and dust emission is the ideal probe of the global disc properties, like size, mass and radial distribution of disc material, because the bulk of the disc mass is located beyond 100 AU from the star, at temperatures of 10-50 K that dominate this part of the spectrum."787 Dise models which include constraints of both dust and molecular gas observations have stressed the importance of analysing the gas and dust components simultaneously. in the context of a ," Disc models which include constraints of both dust and molecular gas observations have stressed the importance of analysing the gas and dust components simultaneously, in the context of a ."788Until recently. observations of rotational transitions of molecules in the submillimetre regime were focused primarily on the low-/ emission from 0Ο. up to the J 23-2 line(22222).," Until recently, observations of rotational transitions of molecules in the submillimetre regime were focused primarily on the $J$ emission from $^{12}$ CO, up to the $J=$ 3-2 line."789" In two of the brightest and most studied sources. Hyaand15.. the observations of higher-/ transitions of ""CO. up to J=6-5 (E, =116 Κ). were compared to the low-J lines. providing estimates of the gas temperature in the intermediate-height molecular layer(?).. crucial ingredients for chemical modelling of disces."," In two of the brightest and most studied sources, and, the observations of $J$ transitions of $^{12}$ CO, up to $=$ 6–5 $E_{\rm k}=$ 116 K), were compared to the $J$ lines, providing estimates of the gas temperature in the intermediate-height molecular layer, crucial ingredients for chemical modelling of discs."790" These single-dish line spectra were fitted using simplistic dise models. deriving a temperature of 20-40 K in the ""CO line emitting layers of LkCa 15. and more than 40 K in TW Hya."," These single-dish line spectra were fitted using simplistic disc models, deriving a temperature of 20-40 K in the $^{12}$ CO line emitting layers of LkCa 15, and more than 40 K in TW Hya."791 analysed submillimetric interferometer observations of TW Hya in the context of a dise structure based on an accretion dise model(?)., analysed submillimetric interferometer observations of TW Hya in the context of a disc structure based on an accretion disc model.792. Based on CO J -6-δ. 3-2 and 2-1 observations. they show that ray heating of the gas is efficient in this source. in addition to the stellar radiation field.," Based on $^{12}$ CO $J=$ 6–5, 3–2 and 2–1 observations, they show that X-ray heating of the gas is efficient in this source, in addition to the stellar radiation field."793" Such diagnosties of gas heating and ionisation improve our understanding of how the gas content evolves in discs,", Such diagnostics of gas heating and ionisation improve our understanding of how the gas content evolves in discs.794 The emerging (sub)millimetre facilities in the Southern hemisphere like the Atacama Pathfinder EXperiment (APEX) and the Australia Telescope Compact Array (ATCA) are opening a window towards the star-forming regions of the Southern sky and are well suited to study circumstellar disc emission., The emerging (sub)millimetre facilities in the Southern hemisphere like the Atacama Pathfinder EXperiment (APEX) and the Australia Telescope Compact Array (ATCA) are opening a window towards the star-forming regions of the Southern sky and are well suited to study circumstellar disc emission.795" These instruments also pave the path for future observations with the Atacama Large Millimetre / Submillimeter Array (ALMA). which will drastically improve our knowledge of dise structure and evolution2).. We use APEX receivers APEX-2a and CHAMP"" to observe the ""CO J 27-6. J 26-5. J 23-2. CO J 23-2 and |C 1] P;—PP, line emission towards the disc around the young intermediate-mass star HD 100546."," These instruments also pave the path for future observations with the Atacama Large Millimetre / Submillimeter Array (ALMA), which will drastically improve our knowledge of disc structure and evolution We use APEX receivers APEX-2a and $^+$ to observe the $^{12}$ CO $J=$ 7–6, $J=$ 6–5, $J=$ 3–2, $^{13}$ CO $J=$ 3–2 and [C I] $^3$ $_2$ $^3$ $_1$ line emission towards the disc around the young intermediate-mass star HD 100546."796 A wealth of observations of dust in this bright disc has motivated us to probe its molecular gas content and kinematics., A wealth of observations of dust in this bright disc has motivated us to probe its molecular gas content and kinematics.797 The chosen transitions are particularly sensitive to the gas in the warm upper layers and kinematics of the outer disc., The chosen transitions are particularly sensitive to the gas in the warm upper layers and kinematics of the outer disc.798 Our millimetre line observations probe the outer radius and inclination., Our millimetre line observations probe the outer radius and inclination.799 The existing. observational constraints on these parameters in the disc around HD 100546 provide an excellent basis for the analysis of our data., The existing observational constraints on these parameters in the disc around HD 100546 provide an excellent basis for the analysis of our data.800 Our observations also provide a bridge toward even higher-J far infrared 'CO transitions to be observed with the Herschel Space Observatory., Our observations also provide a bridge toward even $J$ far infrared $^{12}$ CO transitions to be observed with the Herschel Space Observatory.801 HD 100546 is a young B9V type. 2.5 M. star. classified as a Herbig Be star due to its isolation. infrared excess and silicate emission(??).," HD 100546 is a young B9V type, 2.5 $_{\odot}$ star, classified as a Herbig Be star due to its isolation, infrared excess and silicate emission."802. With a distance of 10346 pe. measured by Hipparcos. this is one of the nearest Herbig Ae/Be stars.," With a distance of $\pm$ 6 pc, measured by Hipparcos, this is one of the nearest Herbig Ae/Be stars."803 The age of the star is estimated to be greater than 10 Myr (?)., The age of the star is estimated to be greater than 10 Myr .804. This makes the presence of circumstellar material intriguing. considering that dises are found to dissipate within 10 Myr in. most young stars 2?2)..Based on SED modelling. postulate the presence of an inner hole in," This makes the presence of circumstellar material intriguing, considering that discs are found to dissipate within 10 Myr in most young stars .Based on SED modelling, postulate the presence of an inner hole in"805"with a given ""standard. deviation"" will not be exactly that appropriate to a gaussian.",with a given `standard deviation' will not be exactly that appropriate to a gaussian.806 Nonetheless. the error. distributions measured on the scales of relevance are close to gaussian. as can be seen in figures ?? and ?? below.," Nonetheless, the error distributions measured on the scales of relevance are close to gaussian, as can be seen in figures \ref{pic4} and \ref{pic5}807 below."808 We also adopt. very conservative confidence limits., We also adopt very conservative confidence limits.809 In view of the (null) results below. and the rather short timescale on which I8O data reduction and calibration is evolving. more complex analyses are unjustified.," In view of the (null) results below, and the rather short timescale on which ISO data reduction and calibration is evolving, more complex analyses are unjustified."810 Finally. the photometric measurements. with associated error bars. are tested against a dark halo model bx the minimisation of X7 in à two-parameter model.," Finally, the photometric measurements, with associated error bars, are tested against a dark halo model by the minimisation of $\chi^2$ in a two-parameter model."811 Shown in figure ?? are the galaxy images. with the regions excised indicated. and some example annuli.," Shown in figure \ref{excisions} are the galaxy images, with the regions excised indicated, and some example annuli."812 We analysed above the galaxy rotation curves. to determine the two convenient. parameters py and £2. describing the rotation curve. and hence the clark halo density.," We analysed above the galaxy rotation curves, to determine the two convenient parameters $\rho_0$ and $R_c$ describing the rotation curve, and hence the dark halo density."813 We additionally have the redshift distance d lor cach galaxy., We additionally have the redshift distance $d$ for each galaxy.814 We noted above that the projected Luminosity distribution expected from a simple halo model should. follow where py can be taken to be an observed Luminosity density without loss of generality., We noted above that the projected luminosity distribution expected from a simple halo model should follow where $\rho_0$ can be taken to be an observed luminosity density without loss of generality.815 Before fitting this functional form to our data. however. we must consider one more practicality.," Before fitting this functional form to our data, however, we must consider one more practicality."816 The offset between the background of the control field and of the target field need not be perfectly zero due to gradients in zodiacal light and galactic emission. and any uncorrected drifts in svstenmi sensitivity.," The offset between the background of the control field and of the target field need not be perfectly zero due to gradients in zodiacal light and galactic emission, and any uncorrected drifts in system sensitivity."817 Any such olfset is unimportant. since it corresponds to a Hat zero point olfset. and not a gradient centred on the target galaxy.," Any such offset is unimportant, since it corresponds to a flat zero point offset, and not a gradient centred on the target galaxy."818 Adelitionally. most olfset fields overlap the galaxy. field. providing further independent checks that we are not ignoring dark halos with very flat. central Iuminosity profiles.," Additionally, most offset fields overlap the galaxy field, providing further independent checks that we are not ignoring dark halos with very flat central luminosity profiles."819 The standard procedure in Lt photometry is to subtract (beamswiteh’) the olfset field from the target field. providing a notionallv zero background dataset.," The standard procedure in IR photometry is to subtract (`beamswitch') the offset field from the target field, providing a notionally zero background dataset."820 Llowever. it is more reliable to fit the relevant background. value directly. for each field.," However, it is more reliable to fit the relevant background value directly, for each field."821 Thus. we require an additional parameter. e. to handle this zero point.," Thus, we require an additional parameter, $\sigma_0$, to handle this zero point."822" Phat is. we fit a model of the form where 6 is given in areseconds. and 0, is the angular equivalent. of the ‘core radius fee (ie. Gy =Re fd)."," That is, we fit a model of the form where $\theta$ is given in arcseconds, and $\theta_0$ is the angular equivalent of the `core radius' $R_C$ (i.e. $\theta_0$ $R_C / d$ )."823 The essential astrophysics is now quantified in the parameter o. which is the normalisation of any luminosity associated: with a dark miattor-like density. profile.," The essential astrophysics is now quantified in the parameter $\alpha$, which is the normalisation of any luminosity associated with a dark matter-like density profile."824" The fits of this model to the data are presented in figures ?? and ?7.. while the derived parameters and their uncertainties are presented in table οτι,"," The fits of this model to the data are presented in figures \ref{pic4}825 and \ref{pic5}, while the derived parameters and their uncertainties are presented in table \ref{limitstable}."826 Ehe zero point flux levels are those presented in table 27. above. where they are seen to be in agreement with DIRBE data.," The zero point flux levels are those presented in table \ref{backlevels} above, where they are seen to be in agreement with DIRBE data."827 Iteassuringlv. the best fit model parameters. in the minimised X72 sense. are given for four out," Reassuringly, the best fit model parameters, in the minimised $\chi^2$ sense, are given for four out"828mag fainter.,mag fainter.829 The wings of both neighboring stars encroach into the photometric aperture., The wings of both neighboring stars encroach into the photometric aperture.830 The Keck image of KOI 877 shows no neighboring stars down to 20th mag., The Keck image of KOI 877 shows no neighboring stars down to 20th mag.831 We have not examined these neighboring stars to determine if they are eclipsing binaries., We have not examined these neighboring stars to determine if they are eclipsing binaries.832" However, the centroid statistics given above indicate that, should they prove eclipsing binaries, they are unlikely to be the cause of the observed transit signature."," However, the centroid statistics given above indicate that, should they prove eclipsing binaries, they are unlikely to be the cause of the observed transit signature."833" For KOI 191, the Keck image reveals a neighboring star located 1.5 arcsec east of the main star and 2.6 mag fainter."," For KOI 191, the Keck image reveals a neighboring star located 1.5 arcsec east of the main star and 2.6 mag fainter."834" As stated above, KOI 191 exhibits some correlation in the rain plots which indicates a crowded field."," As stated above, KOI 191 exhibits some correlation in the rain plots which indicates a crowded field."835" The light curve from KOI 191 also shows an additional, periodic transit feature."," The light curve from KOI 191 also shows an additional, periodic transit feature."836 The ephemeris of this feature is Τε—2454900BJD)=65.6589+Ex0.7086 Days., The ephemeris of this feature is $T_{\text{c}} - 2454900 BJD) = 65.6589 + E \times 0.7086$ Days.837 Its V-shaped transit shape (0.2 mmag depth) and its two-hour duration indicate that a faint eclipsing binary is also blended with the target., Its V-shaped transit shape (0.2 mmag depth) and its two-hour duration indicate that a faint eclipsing binary is also blended with the target.838" False positive scenarios were investigated for the five systems by exploring the possibility that the photometry is the result of contamination of the light of the candidate by an eclipsing binary along the same line of sight, a “blend”."," False positive scenarios were investigated for the five systems by exploring the possibility that the photometry is the result of contamination of the light of the candidate by an eclipsing binary along the same line of sight, a “blend”."839" Given that the centroid motion statistics discussed above rule out a large fraction of the background contaminants, we focused here on hierarchical triple systems in which the candidate and the binary are at the same distance."," Given that the centroid motion statistics discussed above rule out a large fraction of the background contaminants, we focused here on hierarchical triple systems in which the candidate and the binary are at the same distance."840 Angular separations in these cases would usually be too small to generate significant centroid motion., Angular separations in these cases would usually be too small to generate significant centroid motion.841 We modeled the photometry of each candidate assuming it is the result of the brightness variations of an eclipsing binary being attenuated by the (typically) brighter candidate star., We modeled the photometry of each candidate assuming it is the result of the brightness variations of an eclipsing binary being attenuated by the (typically) brighter candidate star.842" For KOIs with two or more signals in the light curve, we modeled the light curves at each period separately and accounted for possible blends at the other period(s) by incorporating extra dilution consistent with those other stars."," For KOIs with two or more signals in the light curve, we modeled the light curves at each period separately and accounted for possible blends at the other period(s) by incorporating extra dilution consistent with those other stars."843" In these cases, whether the blended eclipsing systems at each period are related or not (ie., in a hierarchical quadruple system, for candidates with two signals) is immaterial for the purposes of modeling the light curves."," In these cases, whether the blended eclipsing systems at each period are related or not (i.e., in a hierarchical quadruple system, for candidates with two signals) is immaterial for the purposes of modeling the light curves."844" The objects composing the binary are referred to as the “secondary” and “tertiary”, and the candidate is the primary""."," The objects composing the binary are referred to as the “secondary” and “tertiary”, and the candidate is the “primary”."845 The procedure, The procedure846observations presented by Miharaetal.(1991) where the unabsorbed component was clearly required.,observations presented by \citey{Mihara91.1} where the unabsorbed component was clearly required.847 As proposed by Chotetal.(1994) from the analysis of the pulsed fraction of the lighteurve. the unabsorbed component is probably due to scattering of radiation in an extended hot electron corona into the line of sight.," As proposed by \citey{Choi94.1}848 from the analysis of the pulsed fraction of the lightcurve, the unabsorbed component is probably due to scattering of radiation in an extended hot electron corona into the line of sight."849 These authors also show that the interpretation of Ushimaruetal.(1989).. who attributed the unabsorbed component to a leaky cold absorber. does not hold.," These authors also show that the interpretation of \citey{Ushimaru89}, who attributed the unabsorbed component to a leaky cold absorber, does not hold."850 The anticorrelation between Jy; and Vy found by Leahyal.(1994) and Leahy(1997) in observations has been interpreted by these authors as evidence that the obscuring material partially also obscures the extended corona., The anticorrelation between $I_\mathrm{U}$ and $N_\mathrm{H}$ found by \citey{Leahy94.1} and \citey{Leahy97.1} in observations has been interpreted by these authors as evidence that the obscuring material partially also obscures the extended corona.851 Thus. the geometric covering fraction of the obscuring material is assumed to be quite high during episodes of large Vy.," Thus, the geometric covering fraction of the obscuring material is assumed to be quite high during episodes of large $N_\mathrm{H}$ ."852 Our measured values of J) exhibit some variability (Fig., Our measured values of $I_\mathrm{U}$ exhibit some variability (Fig.853 dcc). but there is no systematic correlation between Zi: and Ny. neither has such a correlationbeen seen by Choretal.(1994)...," \ref{fig:timeevol}c c), but there is no systematic correlation between $I_\mathrm{U}$ and $N_\mathrm{H}$, neither has such a correlationbeen seen by \citey{Choi94.1}."854 A possible interpretation for this discrepancy ts that Leahyetal.(1994) and Leahy(1997) did not include Thomson scattering in their fits., A possible interpretation for this discrepancy is that \citey{Leahy94.1} and \citey{Leahy97.1} did not include Thomson scattering in their fits.855 Indeed. when setting &=0 in Eq. (1))," Indeed, when setting $k=0$ in Eq. \ref{eq:pcovmod}) )"856 and fitting the RXTE data with both normalizations as free parameters. i is much more vartable and 74 appears to be correlated with Αμ.," and fitting the RXTE data with both normalizations as free parameters, $I_\mathrm{U}$ is much more variable and $I_\mathrm{A}$ appears to be correlated with $N_\mathrm{H}$ ."857 In addition. it is generally difficult to distinguish between the absorbed and the unabsorbed component for low values of Αι such that 7i: and 74 get easily confused by the fitting routine.," In addition, it is generally difficult to distinguish between the absorbed and the unabsorbed component for low values of $N_\mathrm{H}$ such that $I_\mathrm{U}$ and $I_\mathrm{A}$ get easily confused by the fitting routine."858 We also tried fitting the data with ü spectral model in which 7: was held fixed at 0.005phem?sΚον|. the average value of the fits of Sect. ??..," We also tried fitting the data with a spectral model in which $I_\mathrm{U}$ was held fixed at $0.005\,\mbox{ph}\,\mbox{cm}^{-2}\,\mbox{s}^{-1}\,\mbox{keV}^{-1}$, the average value of the fits of Sect. \ref{subsec:specfits}."859" The resulting 42, values from this fit were comparable to those of the fits presented in Sect.", The resulting $\chi^2_\mathrm{red}$ values from this fit were comparable to those of the fits presented in Sect.860" 2? since the variations in. 7p in the latter fit are small enough to be compensated by slight changes in AQ, in the former.", \ref{subsec:specfits} since the variations in $I_\mathrm{U}$ in the latter fit are small enough to be compensated by slight changes in $N_\mathrm{H}$ in the former.861 Furthermore. due to the lower temporal resolution of the previous observations. small variations of Ny could not be resolved in these data.," Furthermore, due to the lower temporal resolution of the previous observations, small variations of $N_\mathrm{H}$ could not be resolved in these data."862 It has been pointed out by Parmaret (1986).. that this effect might result in a large uncertainty in the determination of 7i.," It has been pointed out by \citey{Parmar86.1}, that this effect might result in a large uncertainty in the determination of $I_\mathrm{U}$."863 As is shown by our fits to structures in the lighteurve (Fig. 8).," As is shown by our fits to structures in the lightcurve (Fig. \ref{fig:lcfit}) ),"864 we are able to resolve and identify individual structures with à temporal resolution of about one minute., we are able to resolve and identify individual structures with a temporal resolution of about one minute.865 Thus we are confident that the investigation presented here is unaffected by these problems., Thus we are confident that the investigation presented here is unaffected by these problems.866 In our analysis of Sect., In our analysis of Sect.867 ??. and ?? we assumed that NoNy=1.21. re. the value appropriate for material of solar composition.," \ref{subsec:specfits} and \ref{subsec:colordiag}868 we assumed that $N_\mathrm{e}/N_\mathrm{H}=1.21$, i.e., the value appropriate for material of solar composition."869 Previous investigations of dipping sources. however. hinted at non-solar abundances in most of these systems (Reynolds&Parmar1995.. Whiteetal. 1905).," Previous investigations of dipping sources, however, hinted at non-solar abundances in most of these systems \cite{Reynolds95.1}, \cite{White}) )."870 A direct measurement of the abundance ought to be possible byfitting the RXTE data with a partial covering model in which NandNy are both free parameters., A direct measurement of the abundance ought to be possible byfitting the RXTE data with a partial covering model in which $N_\mathrm{e}$ and$N_\mathrm{H}$ are both free parameters.871 We find that the average No/Ny= 1.5., We find that the average $N_\mathrm{e}/N_\mathrm{H}=1.5$ .872 Thus. our data could be interpreted as pointing," Thus, our data could be interpreted as pointing"873"correlation function C),(@) maximizes the detection significance S.",correlation function $C_{\rm m}(\theta)$ maximizes the detection significance $S$.874" Then we use this C,(6) to calculate the statistic y.", Then we use this $C_{\rm m}(\theta)$ to calculate the statistic $\gamma$.875" We repeat this, recording the values of y, to establish the statistical distribution of 7 values."," We repeat this, recording the values of $\gamma$, to establish the statistical distribution of $\gamma$ values."876 We take the value of the threshold γαι to be that for which there is less than a 0.1% chance of getting y>7tn for the case of mg=0., We take the value of the threshold $\gamma_{\rm th}$ to be that for which there is less than a $0.1\%$ chance of getting $\gamma>\gamma_{\rm th}$ for the case of $m_{\rm g}=0$.877" After we determined the 7, a fourth Monte-Carlo simulation is used to calculated Py."," After we determined the $\gamma_{\rm th}$, a fourth Monte-Carlo simulation is used to calculated $P_{\rm d}$."878" This simulation is very similar to the previous one, except that a GW background is generated."," This simulation is very similar to the previous one, except that a GW background is generated."879 We repeat the simulation 10? times and take as Py the probability of getting y>., We repeat the simulation $10^{3}$ times and take as $P_{\rm d}$ the probability of getting $\gamma>\gamma_{\rm th}$.880" We summarize the results of the simulation in Fig. 4,,"," We summarize the results of the simulation in Fig. \ref{fig:decprob},"881" which are gray-scale contour plots for the detection rate P4 for different scenarios of using 60, 100, and 300 pulsars respectively."," which are gray-scale contour plots for the detection rate $P_{\rm d}$ for different scenarios of using 60, 100, and 300 pulsars respectively."882 The corresponding parameters are given in the legend of each panel., The corresponding parameters are given in the legend of each panel.883" Intuitively, the necessary conditions for a positive detection of a graviton mass should be first, that the GW is strong enough for the GW to be detected, and second, that the graviton mass is large enough to change the shape of correlation function."," Intuitively, the necessary conditions for a positive detection of a graviton mass should be first, that the GW is strong enough for the GW to be detected, and second, that the graviton mass is large enough to change the shape of correlation function."884" This intuition is confirmed by our simulations, which show that the high detection rate concentrates in the upper right corner of each panel, where both graviton mass and GW amplitude are large enough."," This intuition is confirmed by our simulations, which show that the high detection rate concentrates in the upper right corner of each panel, where both graviton mass and GW amplitude are large enough."885 From Fig., From Fig.886 4 we can also see that we need at least 60 pulsars to be able to tell the difference between a massive GW background and a massless one., \ref{fig:decprob} we can also see that we need at least 60 pulsars to be able to tell the difference between a massive GW background and a massless one.887 For 5-year observations of 100 pulsars we can start to detect a graviton heavier than 2.5x10733 eV and we can achieve a limit of mg=10:33 eV by using 5-year observations of 300 pulsars.," For 5-year observations of 100 pulsars we can start to detect a graviton heavier than $2.5\times10^{-22}$ eV and we can achieve a limit of $m_{\rm 888g}=10^{-22}$ eV by using 5-year observations of 300 pulsars."889 We can achieve levels of 10:33 eV and 5x1077? eV in 10-year observations using 100 and 300 puslars respectively., We can achieve levels of $10^{-22}$ eV and $5\times 10^{-23}$ eV in 10-year observations using 100 and 300 puslars respectively.890" We also note that there is a positive correlation between the minimal detectable graviton mass and the GW background amplitude, as shown by the leftwards lead edge of the contours, in other words, we need a larger GW amplitude such that the GW background can be detected, if gravitons are more massive."," We also note that there is a positive correlation between the minimal detectable graviton mass and the GW background amplitude, as shown by the leftwards lead edge of the contours, in other words, we need a larger GW amplitude such that the GW background can be detected, if gravitons are more massive."891" As discussed above, this correlation is due to the reduction of pulsar timing residuals due to massive graviton."," As discussed above, this correlation is due to the reduction of pulsar timing residuals due to massive graviton."892de-noised liebt curve.,de-noised light curve.893 The removed scales are the first scale of details. that correspond with high frequencies. as well as the last 6. 7 or 8 scales of approximations that correspond with low frequencies.," The removed scales are the first scale of details, that correspond with high frequencies, as well as the last 6, 7 or 8 scales of approximations that correspond with low frequencies."894 Tow many ApC scales are multiplied by zero depends on the signal., How many ApC scales are multiplied by zero depends on the signal.895 The selection is performed automatically by measuring the dispersion of the lieht curves: the higher the dispersion the stronger the filter. that is. inore scales are removed.," The selection is performed automatically by measuring the dispersion of the light curves; the higher the dispersion the stronger the filter, that is, more scales are removed."896 The fiction used as a imnother wavelet was a Daubechies order 21 (Daubechies 19923) ., The function used as a mother wavelet was a Daubechies order 24 (Daubechies \cite{daub}) ).897 Although many different functious that produce different cocficieuts cau be used as a mother wavelet. for our filtering process the final result is almost independent of the used function.," Although many different functions that produce different coefficients can be used as a mother wavelet, for our filtering process the final result is almost independent of the used function."898 The result of this filter is shown in Fie., The result of this filter is shown in Fig.899 2 (bottom) for the helt curve 168., 3 (bottom) for the light curve 168.900 TRUFAS starts with a Continuous Wavelet Transtormation (οΝΤ) (Torrence and Compo 1998)}) of the filtered data. which is used to select the part of the signal where the transit search will be done.," TRUFAS starts with a Continuous Wavelet Transformation (CWT) (Torrence and Compo \cite{torrence}) ) of the filtered data, which is used to select the part of the signal where the transit search will be done."901 The continuous wavelet transform of a function fy) is defined. by Ou “nother” wavelet was a Paul function (Torrence aud Compo 1998)]: and the scaled wavelet is: where gp is the time in this case. s the dilation paraiucter used to change the scale. aud η the translation paramcter used to slide in fine.," The continuous wavelet transform of a function $f({\eta})$ is defined by: Our “mother” wavelet was a Paul function (Torrence and Compo \cite{torrence}) ): and the scaled wavelet is: where $\eta$ is the time in this case, $s$ the dilation parameter used to change the scale, and $n$ the translation parameter used to slide in time."902 The factor of s1? is a normalisation factor to keep the total enerew of the scaled wavelet coustaut., The factor of $s^{-1/2}$ is a normalisation factor to keep the total energy of the scaled wavelet constant.903 Therefore. the CWT maps the signal iu a two dimension function in a time-scale space.," Therefore, the CWT maps the signal in a two dimension function in a time-scale space."904 The decomposition was made in 55 scales., The decomposition was made in 55 scales.905 The Paul fiuction order l was selected because its shape is similar to the feature we are looking for., The Paul function order 1 was selected because its shape is similar to the feature we are looking for.906 In fact. the continuous wavelet transform is just a correlation between he function. the helt curve in our case. aud the scaled aud shifted wavelet function.," In fact, the continuous wavelet transform is just a correlation between the function, the light curve in our case, and the scaled and shifted wavelet function."907 The higher the correlation. the Πσ]ο the cocficieut of the CWT.," The higher the correlation, the higher the coefficient of the CWT."908 From the 55 used scales. see Fie. L.," From the 55 used scales, see Fig. \ref{fig55scales},"909 zu automatic selection of the best scale was xxformed., an automatic selection of the best scale was performed.910 The selectiou of he scale was done following he double criteria that the best scale is the one with üeher cocficicuts when these high coefficients are present iu more scales., The selection of the scale was done following the double criteria that the best scale is the one with higher coefficients when these high coefficients are present in more scales.911 A scale may be considered to correspond o a transit with a eiven duration if the width of the central. Gaussian-like. part of the scale has a duration similar to that transit.," A scale may be considered to correspond to a transit with a given duration if the width of the central, Gaussian-like, part of the scale has a duration similar to that transit."912 When he automatic scale-selection eave results outside the range correspondiug to transits. between 2.3 and 9.5 hours (iu the DTI set of test leh curves. one nav only expect durations in that range). scales that correspond to transit durations of 5.7 h were used. being a good compromise for the expected lcugh of transits.," When the automatic scale-selection gave results outside the range corresponding to transits, between 2.3 and 9.5 hours (in the BT1 set of test light curves, one may only expect durations in that range), scales that correspond to transit durations of 5.7 h were used, being a good compromise for the expected lenght of transits."913 Such failures of the automatic scale sclection happen m cases of transits of very low S/N. The use of scales significantly deviating from the expected transi duratiou could introduce problems iu the next step of the algoritlan: when the scales are too narrow. peaks related to lugh frequency noise could produce false detections in the transit-scarch algorithia: on the other haud. scales that are too wide produce signals that are also too wide to he a suitable input to the peak-searching algoritlan that is used in the next step.," Such failures of the automatic scale selection happen in cases of transits of very low S/N. The use of scales significantly deviating from the expected transit duration could introduce problems in the next step of the algorithm: when the scales are too narrow, peaks related to high frequency noise could produce false detections in the transit-search algorithm; on the other hand, scales that are too wide produce signals that are also too wide to be a suitable input to the peak-searching algorithm that is used in the next step."914 Fie., Fig.915" 5 shows two exanples o: selected scales,", \ref{figselscale} shows two examples of selected scales.916 As it was pointed out previously. the selected scale is a one dimensional fiction of the correlation coeficicut vetween the light curve aud the wavelet trausforii. versus ine.," As it was pointed out previously, the selected scale is a one dimensional function of the correlation coefficient between the light curve and the wavelet transform, versus time."917 When a planetary transi with high S/N is prescut in the data (Fig. 5..," When a planetary transit with high S/N is present in the data (Fig. \ref{figselscale},"918 op). it appears as equally spaced oeaks in the selected scale.," top), it appears as equally spaced peaks in the selected scale."919 Transits whose amplitude is colparable to the noise will lead to a set of peaks (Fig. 5..," Transits whose amplitude is comparable to the noise will lead to a set of peaks (Fig. \ref{figselscale},"920 tton). sole from the transits and some from nolse. without apparent periodicity.," bottom), some from the transits and some from noise, without apparent periodicity."921 The search for periodicities among these peaks is doue iu the next step. using the vower (the square of the signal) of the selected scale.," The search for periodicities among these peaks is done in the next step, using the power (the square of the signal) of the selected scale."922 If a Πο curve contains transits in it. there will be a set of peaks regularly spaced in time by an amount that coincides with the period of the planet.," If a light curve contains transits in it, there will be a set of peaks regularly spaced in time by an amount that coincides with the period of the planet."923 It is this feature or pattern that defines the signature of the transit., It is this feature or pattern that defines the signature of the transit.924 The level of the amplitude of the transit relative to the noise, The level of the amplitude of the transit relative to the noise925predict the halo correlation function. even for Gaussian nmoclels.,predict the halo correlation function even for Gaussian models.926 To summarize. cluster correlation length observations place strong cosmology independent constraints on. the matter power spectrum in the universe. constraints which [uture surveys such as SDSS will allow us to fully exploit.," To summarize, cluster correlation length observations place strong cosmology independent constraints on the matter power spectrum in the universe, constraints which future surveys such as SDSS will allow us to fully exploit."927 ] would like to thank Marc Davis. Lvic Gawiser and Joseph Silk for helpful ancl stimulating discussions.," I would like to thank Marc Davis, Eric Gawiser and Joseph Silk for helpful and stimulating discussions."928 E. would. also like to thank Fabio Governato and LH. J. Mo for their swift answers to mv questions. and. Pedro Ferreira for reading a draft of the manuscript.," I would also like to thank Fabio Governato and H. J. Mo for their swift answers to my questions, and Pedro Ferreira for reading a draft of the manuscript."929 This work has been supported. in part by erants from the NSE. including grant 9617168.," This work has been supported in part by grants from the NSF, including grant 9617168."930Assuming a ean X-ray count-rate-to-flux conversion factor of 1.1. ς 1t ore + 7. which we derive frou A-vav spectral fits of the TTS in Sect. 6..,"Assuming a mean X-ray count-rate-to-flux conversion factor of 1.1 $\times$ $^{-11}$ erg $^{-1}$ $^{-2}$, which we derive from X-ray spectral fits of the TTS in Sect. \ref{xlf},"931 if the cloud is at a distance of 65 pc. the linitius huninosities of the observations are 1.7.«1029 ere + aud 7.2«107* eres + for the RASS audROSAT poiuted observatious. respectively.," if the cloud is at a distance of 65 pc, the limiting luminosities of the observations are $1.7 \times 10^{29}$ erg $^{-1}$ and $7.2 \times 10^{27}$ erg $^{-1}$ for the RASS and pointed observations, respectively."932 Therefore. these observations are sufiicicut to detect most of the WITS iu the cloud since the threshold is below the N-rav füutest stars in the WITS N-rav hDmuuinositv function (e.g... Neuliumser et al.," Therefore, these observations are sufficient to detect most of the WTTS in the cloud since the threshold is below the X-ray faintest stars in the WTTS X-ray luminosity function (e.g., Neuhäuuser et al."933 1995)., 1995).934 Although the RASS observation. of MDBMI2 in not sensitive chough to detect all of the CTTS in the cloud. the objective prisii survey by Stephenson (1986) identified all of he Πα cmission sources in this region down to a visual maeuitude threshold of ~ |3.5.," Although the RASS observation of MBM12 in not sensitive enough to detect all of the CTTS in the cloud, the objective prism survey by Stephenson (1986) identified all of the $\alpha$ emission sources in this region down to a visual magnitude threshold of $\sim$ 13.5."935 Since this lamitine magnitude corresponds to the carly AL spectral types in MBAIL2. the current population of TTS in MDMI2 presented in this paper should be complete for all carlicr spectral types.," Since this limiting magnitude corresponds to the early M spectral types in MBM12, the current population of TTS in MBM12 presented in this paper should be complete for all earlier spectral types."936 Since MDMAI2 is at relatively ligh galactic latitude. many of the Sl N-raw ποος WC ideutified are extragalactic.," Since MBM12 is at relatively high galactic latitude, many of the 81 X-ray sources we identified are extragalactic."937 Therefore. we used the X-ray to optical flux ratios (sce Table 1) to remove extragalactic sources from our List of candidates (cf.," Therefore, we used the X-ray to optical flux ratios (see Table 4) to remove extragalactic sources from our list of candidates (cf."938 Tearty et al., Hearty et al.939 1999)., 1999).940 All sources which have loe( f/f} 0.0 are considered to be extragalactic and those with Ίου) < 0.0 are considered to be stellar objects. some of which could be PAIS.," All sources which have $f_{\rm x}/f_{\rm v}$ ) $>$ 0.0 are considered to be extragalactic and those with $f_{\rm x}/f_{\rm v}$ ) $<$ 0.0 are considered to be stellar objects, some of which could be PMS."941 We also searched the literature to remove cataloged 1on-PAIS stars from our list of cauclicates., We also searched the literature to remove cataloged non-PMS stars from our list of candidates.942 Finally we were eft with a list of N-rav sources ideutified in the RASS andROSAT pointed observations of the cloud which ave stellar optical counterparts that may be PAIS stars., Finally we were left with a list of X-ray sources identified in the RASS and pointed observations of the cloud which have stellar optical counterparts that may be PMS stars.943" However, nauv of these stars may be other types of X-rav active stars (ce... RS CVn aud dMoe stars) aud nearby hain sequence stars (which may not be iutrinsicallv X-rav bright. but are near enough so that their X-ray flux is aree) that are more difficult to separate frou: PAIS stars w A-rav observations alone."," However, many of these stars may be other types of X-ray active stars (e.g., RS CVn and dMe stars) and nearby main sequence stars (which may not be intrinsically X-ray bright, but are near enough so that their X-ray flux is large) that are more difficult to separate from PMS stars by X-ray observations alone."944 Therefore. follow-up spectral observations are necessary to identify which X-ray sources are T Τα stars.," Therefore, follow-up spectral observations are necessary to identify which X-ray sources are T Tauri stars."945 Iu order to complete the census of the TTS population of AIDMI2 we require follow-up observations., In order to complete the census of the TTS population of MBM12 we require follow-up observations.946 Since lithium is burned quickly iu convective stars. a measuremeut of W(Li) alone with a knowledge of the spectral tvpe of a star can be a reliable indicator of voutl.," Since lithium is burned quickly in convective stars, a measurement of W(Li) along with a knowledge of the spectral type of a star can be a reliable indicator of youth."947 Therefore we obtaiuec broad-band. low-resolution. optical spectra of the N-rav chuitting TTS caudidates to determine spectral types auk measure the equivalent width of the Πα cmiussion aud GTÜUS aabsorption lues.," Therefore we obtained broad-band, low-resolution, optical spectra of the X-ray emitting TTS candidates to determine spectral types and measure the equivalent width of the $\alpha$ emission and 6708 absorption lines."948 The spectra were obtained frou: October 911. 1998 with the Calar Alto Faint Object Spectrograph (CAFOS) at the 2.220 telescope at Calar Alto. Spain.," The spectra were obtained from October 9–11, 1998 with the Calar Alto Faint Object Spectrograph (CAFOS) at the 2.2-m telescope at Calar Alto, Spain."949 The [jan pixels of the SITe-ld «2018 chip with the C-100, The $\mu$ m pixels of the SITe-1d $\times$ 2048 chip with the G-100950The absence of evidence of any relatively wide pre-magnetic CVs in some 1.253 WDM pairs that have so far been studied sugeests that a MCN must be born preferentially in a semidetached state. as an intermediate polar or AM Ler svstem. or sullicientIv close to contact with its Roche lobe for the HEMWD to capture the wind of its companion so that it appears as à low accretion rate polar.,"The absence of evidence of any relatively wide pre-magnetic CVs in some $1{,}253$ WD+M pairs that have so far been studied suggests that a MCV must be born preferentially in a semidetached state, as an intermediate polar or AM Her system, or sufficiently close to contact with its Roche lobe for the HFMWD to capture the wind of its companion so that it appears as a low accretion rate polar."951 A common envelope magnetic dvnamo leads naturally to a scenario in which this occurs., A common envelope magnetic dynamo leads naturally to a scenario in which this occurs.952 In this model systems that emerge more tightly bound following the CL phase are the systems with the stronger magnetic fields., In this model systems that emerge more tightly bound following the CE phase are the systems with the stronger magnetic fields.953 A consequence of this is that the AICVs might. be expected to have intrinsically hotter white dwarf primaries as the systems first come into contact., A consequence of this is that the MCVs might be expected to have intrinsically hotter white dwarf primaries as the systems first come into contact.954 However. white dwarfs cool in about 10! vr to ellective temperatures. of about 15.000Ix. typical of currently observed ALCWs.," However white dwarfs cool in about $10^7\,$ yr to effective temperatures of about $15{,}000\,$ K typical of currently observed MCVs."955 Ata »eriod. Of 2 hr the gravitational radiation timescale is about 310 vr.," At a period of $2\,$ hr the gravitational radiation timescale is about $3\times 10^9\,$ yr."956 Indeed it is long enough to cool to the observed ellective temperatures of the LARPs., Indeed it is long enough to cool to the observed effective temperatures of the LARPs.957 Phe systematically uigher cllective temperatures seen in the white dwarfs in nonmagnetic CVs are attributed to the relative importance of compressional heating in these svstems owing to their ügher mean acerction rates (Sion2004)., The systematically higher effective temperatures seen in the white dwarfs in nonmagnetic CVs are attributed to the relative importance of compressional heating in these systems owing to their higher mean accretion rates \citep{sion2004}.958. We note that there are some dillerences. in. the sropertics of the isolated HENWDs. anc the AICVs which have previously been attributed. to their cillerent origins., We note that there are some differences in the properties of the isolated HFMWDs and the MCVs which have previously been attributed to their different origins.959 First the incidence of magnetism in CVs (about 25 per cent) is sienificanthy higher than in the isolated white cdwarfs (about LOper cent).," First the incidence of magnetism in CVs (about $25\,$ per cent) is significantly higher than in the isolated white dwarfs (about $10\,$ per cent)."960 This has also been attributed. to. possible selection cllects (Warner1995) though no detailed studies have been carried. out to test this hypothesis., This has also been attributed to possible selection effects \citep{warner1995} though no detailed studies have been carried out to test this hypothesis.961 Secondly the polars are deficient of white dwarfs with fields in excess of 100 compared. with the isolated LENNDs.," Secondly the polars are deficient of white dwarfs with fields in excess of $10^8\,$ G compared with the isolated HFMWDs."962 Our hypothesis accounts for both these dilferences straüghtForwardlv., Our hypothesis accounts for both these differences straightforwardly.963 Phe highest fields are expected in the CEs in which the cores are closest together., The highest fields are expected in the CEs in which the cores are closest together.964 Thus the isolated HEPMWDs can generally have much higher. fields than the MCVs because the differential rotation in their progenitor CI5 can be much greater., Thus the isolated HFMWDs can generally have much higher fields than the MCVs because the differential rotation in their progenitor CE can be much greater.965 Lt is greatest just before mereing., It is greatest just before merging.966 The MCVs then fall in a small range of core separations bordering on the svstems that merge., The MCVs then fall in a small range of core separations bordering on the systems that merge.967 Wider core separations and the end of the CIS phase end up as the ore-cataclvsmic non-magnetic variables anc wider systems hat will never interact., Wider core separations and the end of the CE phase end up as the pre-cataclysmic non-magnetic variables and wider systems that will never interact.968 These coupled. with single stars and binary stars that never enter a CL phase make up the ονΠο] white ciwarls in cataclysmic variables. wider binary stars and single isolated white cbiwarfs.," These coupled with single stars and binary stars that never enter a CE phase make up the low-field white dwarfs in cataclysmic variables, wider binary stars and single isolated white dwarfs."969 The distribution of the masses of the isolated LEEMWDs (Lig. 11} , The distribution of the masses of the isolated HFMWDs (Fig. \ref{mass}) )970appears to be mace up of a distribution similar to hat of the low-Lielel cdwarks augmented from about 0.6AL. upwards and especially so at the very. high masses.," appears to be made up of a distribution similar to that of the low-field dwarfs augmented from about $0.6\,M_\odot$ upwards and especially so at the very high masses."971 Under our hypothesis these are simply explained as the result. of a common envelope with two degenerate cores., Under our hypothesis these are simply explained as the result of a common envelope with two degenerate cores.972 “Typically this is the second Cle phase in a system with two stars that can both evolve to become giants., Typically this is the second CE phase in a system with two stars that can both evolve to become giants.973 The first phase leaves a closer. MS|WD svstem., The first phase leaves a closer MS+WD system.974 The second star then evolves and unstable mass transfer leads to the second. CE phase in which the eiant-like envelope surrounds two degenerate cores., The second star then evolves and unstable mass transfer leads to the second CE phase in which the giant-like envelope surrounds two degenerate cores.975 HE the two cores merge to form a massive WD it has a high magnetic field in accordance with our hypothesis., If the two cores merge to form a massive WD it has a high magnetic field in accordance with our hypothesis.976 Lf 10 total mass exceeds the Chandrasekhar limit the cores may undergo accretion induced. collapse to leave a highly magnetic neutron star., If the total mass exceeds the Chandrasekhar limit the cores may undergo accretion induced collapse to leave a highly magnetic neutron star.977 In either case we expect aecretion uring the merging to be fast enough to burn any material non-degenerately to oxygen and neon (Martin.Tout&Lesaf-re2006)., In either case we expect accretion during the merging to be fast enough to burn any material non-degenerately to oxygen and neon \citep{martin2006}.978. Phese stars ought to emerge rapidly spinning but garoulcl also spin down rapidlv by magnetic braking., These stars ought to emerge rapidly spinning but should also spin down rapidly by magnetic braking.979 One μαar. EUVE J 0317-855 (Ferrarioetal.1997).. shows both a ugh spin. 2=12min. and a high mass. M=1.35M...," One star, EUVE J 0317-855 \citep{ferrario97}, shows both a high spin, $P = 12\,$ min, and a high mass, $M=1.35\,M_{\odot}$."980 The magnetic dvnamo model for the origin of fields in MCVs should be contrasted: with the fossil field. model wt has been used to explain the properties of the isolated —FAIWDs (Wickramasinghe&Ferrario2005)., The magnetic dynamo model for the origin of fields in MCVs should be contrasted with the fossil field model that has been used to explain the properties of the isolated HFMWDs \citep{wickramasinghe2005}.981.. Accorcling o this hypothesis the magnetic flux in the core of the star that becomes the white cwarf is closely related. to 16 Dux of its main-sequence. progenitor., According to this hypothesis the magnetic flux in the core of the star that becomes the white dwarf is closely related to the flux of its main-sequence progenitor.982 The mechanism. w which this correspondence occurs has been unclear but recent calculations (Zahn.Brun&Mathis2007) confirm iu a [fossil field in radiative regions could act as a seed icld to generate a strong magnetic field by a dynamo in 1ο convective core that becomes the white cdwarl (ος.Tout.Wickramasinghe&Ferrario 2004).," The mechanism by which this correspondence occurs has been unclear but recent calculations \citep{zahn2007}983 confirm that a fossil field in radiative regions could act as a seed field to generate a strong magnetic field by a dynamo in the convective core that becomes the white dwarf \citep[c.f.][]{tout2004}."984. Xn alternative dvnamo origin in single stars was proposed by Levy&Rose rut it still relies on the intrinsic stellar spin of a single star., An alternative dynamo origin in single stars was proposed by \citet{levy1974} but it still relies on the intrinsic stellar spin of a single star.985 Phe major dillerence in our hypothesis is that. unlike in single star evolution. strong cillerential rotation is anessential characteristic of common envelope evolution.," The major difference in our hypothesis is that, unlike in single star evolution, strong differential rotation is an characteristic of common envelope evolution."986 Our assumption is that any intrinsic fossil field that is present is destroved or at most serves as a seed field for the magnetic dvnamo in the CIS phase., Our assumption is that any intrinsic fossil field that is present is destroyed or at most serves as a seed field for the magnetic dynamo in the CE phase.987 The space densities and observable lifetimes of CVs and white dwarfs are not well known but we can use current estimates to check consisteney with our hypothesis., The space densities and observable lifetimes of CVs and white dwarfs are not well known but we can use current estimates to check consistency with our hypothesis.988 There are 1.12.3107 CVs per cubic parsee (Pretoriusetal.2007) and about one quarter of these have high fields., There are $1.1\pm 2.3\times 10^{-5}$ CVs per cubic parsec \citep{pretorius2007} and about one quarter of these have high fields.989 Phere are 35107 white cwarfs per cubic parsec (Liebertetal.2005) and about one tenth of these are HEPMWDs., There are $3\times 10^{-3}$ white dwarfs per cubic parsec \citep{liebert2005} and about one tenth of these are HFMWDs.990 IH we assume that the observable lifetime of a CV. the time over which its mass-transfer rate is sullicientIy high. is about one tenth that of an isolated white chwarl the time for it to cool below detectable limits. then the birth rate of HEMW is about three times that of CVs.," If we assume that the observable lifetime of a CV, the time over which its mass-transfer rate is sufficiently high, is about one tenth that of an isolated white dwarf, the time for it to cool below detectable limits, then the birth rate of HFMWDs is about three times that of CVs."991 Thus three times as manyDs systems entering a Cl phase should. end up merging. as emerge separated but close enough to become a CV., Thus three times as many systems entering a CE phase should end up merging as emerge separated but close enough to become a CV.992 The fact that no white cwarf with a surface magnetic field over 3BAICG has been found. in a detached. binary system suggests that all such highly maenctic white cwarls have a binary origin.," The fact that no white dwarf with a surface magnetic field over $3\,$ MG has been found in a detached binary system suggests that all such highly magnetic white dwarfs have a binary origin."993 ΙΓ half of the stars in our neighbourhood have a binary companion anc half of these are sulliciently separated not to have interacted then there should. be at least one quarter as many magnetic white dwarfs in wide detached: binary systems as appear as single stars unless their origin depends on binary. interaction., If half of the stars in our neighbourhood have a binary companion and half of these are sufficiently separated not to have interacted then there should be at least one quarter as many magnetic white dwarfs in wide detached binary systems as appear as single stars unless their origin depends on binary interaction.994 Εις is not the case and so argues very strongly against any single star evolutionary origin of ILEMWDs and. very much in favour of a mechanism that relies on binary interaction., This is not the case and so argues very strongly against any single star evolutionary origin of HFMWDs and very much in favour of a mechanism that relies on binary interaction.995 Given the conundrum. relating to the AICVs. that there is an absence of evidence for any pre-ALCVs in some," Given the conundrum, relating to the MCVs, that there is an absence of evidence for any pre-MCVs in some"996large sample of z>4 QSOs [rom the Automatic Plate Measuring survey (Irwin. MeMahon. ΠΠανανά 1991). including the two sources BRI 12020725 al 2=4.7 and BRI 13350415 ab z=4.4.,"large sample of $z \ge 4$ QSOs from the Automatic Plate Measuring survey (Irwin, McMahon, Hazard 1991), including the two sources BRI 1202–0725 at $z = 4.7$ and BRI 1335–0415 at $z = 4.4$."997 Optical spectra of both of these sources show strong associated Lya absorption (Storrie-Lomb:ordi et al., Optical spectra of both of these sources show strong associated $\alpha$ absorption (Storrie-Lombardi et al.998 1996)., 1996).999 Omont οἱ al. (, Omont et al. (10001996) and Ohta et al. (,1996) and Ohta et al. (10011996) detect high order CO line emission from 12020725. as well as thermal continuum emission from warm dust at 1.35mim.,"1996) detect high order CO line emission from 1202–0725, as well as thermal continuum emission from warm dust at 1.35mm."1002" The continuum and line emission consist of a double source withan angular separation of 4"".", The continuum and line emission consist of a double source withan angular separation of $''$.1003 The total flux density of the source at 1.2mm is 1243 mJ. implying a far Dt Iuminosity of Ly=42xLOM L.. assuming a spectral energy. distribution (SED) twpical for an —utva-liminous infrared galaxy (ULIRG: Lygi~ 107). where [μμ is defined as in Helou et al. (," The total flux density of the source at 1.2mm is $12 \pm 3$ mJy, implying a far IR luminosity of $L_{\rm FIR} = 4.2 \times100410^{13}$ $_\odot$, assuming a spectral energy distribution (SED) typical for an ultra-luminous infrared galaxy (ULIRG; $L_{\rm FIR} \sim100510^{12}$ ), where $L_{\rm FIR}$ is defined as in Helou et al. ("10061988) ancl Condon (1992). ie.,"1988) and Condon (1992), ie."1007 the integrated luminosity between rest [rame wavelengths of 42 to 122 yan. The southern source in 12020725 comprises about 65% of the total., the integrated luminosity between rest frame wavelengths of 42 to 122 $\mu$ m. The southern source in 1202–0725 comprises about $\%$ of the total.1008 Observations of this source with ISO at mid- to far-IR wavelengths constrain the dust temperature to be 70 Ix for a dust emissivity index ο=1.5 (Leech. Metcalfe. Altieri 2001).," Observations of this source with ISO at mid- to far-IR wavelengths constrain the dust temperature to be $\simeq 70$ K for a dust emissivity index $\beta = 1.5$ (Leech, Metcalfe, Altieri 2001)."1009 The CO observations of Omont et al. (, The CO observations of Omont et al. (10101996a) aud. Guelin et al. (,1996a) and Guelin et al. (10112001) indicate a difference in the CO line widths for the northern ancl southern components.,2001) indicate a difference in the CO line widths for the northern and southern components.1012 Gaussian fitting to the CO(5-4) line emission profiles bv Omont et al. (, Gaussian fitting to the CO(5-4) line emission profiles by Omont et al. (10131996) results in z=4.6947. line Full Width at Half. Maximum (FWIIM) = 190 kins !. and an integrated line [lux density = I.1z:0.2 Jv lan ! for the southern source. and 2=4.6916. FWIIM = 350 kms |. and an integrated line flux density = 1.320.3 Jv kins | for the northernsource.,"1996) results in $z = 4.6947$, line Full Width at Half Maximum (FWHM) = 190 km $^{-1}$ , and an integrated line flux density = $\pm$ 0.2 Jy km $^{-1}$ for the southern source, and $z = 4.6916$, FWHM = 350 km $^{-1}$ , and an integrated line flux density = $1.3 \pm 0.3$ Jy km $^{-1}$ for the northernsource."1014 This difference, This difference1015"given exposure, but gets uncorrelated rapidly between distinct exposures (as long as they are not taken in immediate succession).","given exposure, but gets uncorrelated rapidly between distinct exposures (as long as they are not taken in immediate succession)."1016 The same holds for some types of instrumental systematics which are not correlated between exposures taken on dillerent nights., The same holds for some types of instrumental systematics which are not correlated between exposures taken on different nights.1017 The systematic errors eliminated by this technique are what we have been calling---- 9$...and dg)u," The systematic errors eliminated by this technique are what we have been calling $\delta \hat\xi_{(\ref{error_atm})}^{(i,i)}$ and $\delta1018\hat\xi_{(\ref{error_a})}^{(i,i)}$."1019" So for these two components of the error, δε=(."," So for these two components of the error, $\delta\hat\xi = 0$."1020" For the three-point (or //—point) correlation function, the same argument holds as long we have at least three (or 1) exposures."," For the three-point (or $n-$ point) correlation function, the same argument holds as long we have at least three (or $n$ ) exposures."1021" With the shapes all taken [rom diflerent exposures, the systematic errors [rom the atmosphere and the coefficient estimates are eliminated, leaving only the errors [rom the PC measurements and the neglected PCs to contribute to the systematic error."," With the shapes all taken from different exposures, the systematic errors from the atmosphere and the coefficient estimates are eliminated, leaving only the errors from the PC measurements and the neglected PCs to contribute to the systematic error."1022" By δές we have been referring to a systematic error, or bias; that is. a change in the expectation value relative to the correct value."," By $\delta\hat\xi$, we have been referring to a systematic error, or bias; that is, a change in the expectation value relative to the correct value."1023 So correlating across multiple exposures results in no systematic error [rom the effects we have numbered 3 and 4..," So correlating across multiple exposures results in no systematic error from the effects we have numbered \ref{error_atm}1024 and \ref{error_a}."1025" However, these errors (all four, actually) also contribute to the statistical error in £, since the variance of£ due to these two errors does not vanish."," However, these errors (all four, actually) also contribute to the statistical error in $\xi$, since the variance of $\xi$ due to these two errors does not vanish."1026" This contribution to the stausucal error has vet to be accurately estimated, but we expect it to be small compared to the sample variance plus intrinsic ellipticity errors Cejntrinsie Irom Equation 4))."," This contribution to the statistical error has yet to be accurately estimated, but we expect it to be small compared to the sample variance plus intrinsic ellipticity errors $\epsilon_{\rm intrinsic}$ from Equation \ref{eqn:psf_error}) )."1027 Wittman(2005) used a set of exposures of a single field imaged with the Subaru telescope to estimate the atmospheric contribution., \citet{Wi05} used a set of exposures of a single field imaged with the Subaru telescope to estimate the atmospheric contribution.1028" This is à concern on arcminute scales or smaller, l'or which the PSF correction may not be accurate even with PCA interpolation if the PSF pattern is non-recurrent."," This is a concern on arcminute scales or smaller, for which the PSF correction may not be accurate even with PCA interpolation if the PSF pattern is non-recurrent."1029 Wittman(2005). finds a contribution to the shear correlation of order 10.? on areminute scales., \citet{Wi05} finds a contribution to the shear correlation of order $10^{-5}$ on arcminute scales.1030" This may be compared to the contribution from intrinsic ellipücities, which is of order 10.7, but scales inversely with the total number of galaxy pairs."," This may be compared to the contribution from intrinsic ellipticities, which is of order $10^{-3}$, but scales inversely with the total number of galaxy pairs."1031" The atmospheric contribution scales inversely with the number of independent coherent patches, which depends on the coherence scale of the atmosphere."," The atmospheric contribution scales inversely with the number of independent coherent patches, which depends on the coherence scale of the atmosphere."1032" Unless this scale is much larger than an arcminute, the atmospheric contribution will be comparatively small."," Unless this scale is much larger than an arcminute, the atmospheric contribution will be comparatively small."1033" The contribution of the other three errors lo the statistical error budget is also likely to be small, but it needs to be estimated for planned surveys."," The contribution of the other three errors to the statistical error budget is also likely to be small, but it needs to be estimated for planned surveys."1034 How many exposures should one take per pointing?, How many exposures should one take per pointing?1035 We have advocated multiple exposures in the discussion above to be able to use galaxies in diflerent exposures to measure shear correlations., We have advocated multiple exposures in the discussion above to be able to use galaxies in different exposures to measure shear correlations.1036" While this eliminates certain systematic errors, by omitting the 7=j terms in the correlation function estimates, we are losing some information."," While this eliminates certain systematic errors, by omitting the $i=j$ terms in the correlation function estimates, we are losing some information."1037" Assume each pair of galaxies which are being used for the two-point correlation function are each observed on Nox)H exposures and have a measurement error, e,,:4,. ON each exposure equal to o\/Nexpfur (80 the measurement error on a stacked image would be c)."," Assume each pair of galaxies which are being used for the two-point correlation function are each observed on $N_{\rm exp}$ exposures and have a measurement error, $\epsilon_{\rm meas}$, on each exposure equal to $\sigma \sqrt{N_{\rm exp}}$ (so the measurement error on a stacked image would be $\sigma$ )."1038" The variance of€ when the 7=j pairs are neglected is l'ound to be: For well measured galaxies, the last term is negligible (assuming .Vexp7 2): however, for faint galaxies,"," The variance of $\xi$ when the $i=j$ pairs are neglected is found to be: For well measured galaxies, the last term is negligible (assuming $N_{\rm exp} \ge 2$ ); however, for faint galaxies,"1039should become weaker at higher frequencies. when the strong pulses iu the leading componcut clearly dominate he energetics.,"should become weaker at higher frequencies, when the strong pulses in the leading component clearly dominate the energetics."1040 That is indeed the case (see Fig. 18)), That is indeed the case (see Fig. \ref{correlc1c2}) )1041 where he flux deusities appear almost anticorrelated at 1500 ΑΠΕ., where the flux densities appear almost anti-correlated at 4850 MHz.1042 This observation also explains the fact uentioned iu Sect., This observation also explains the fact mentioned in Sect.1043 L1 that the correlation cocfitcicut between frequency ors are larecr for cach compoucut separately than for the ull profile., 4.1 that the correlation coefficient between frequency pairs are larger for each component separately than for the full profile.1044 We can study the xoadband. characteristics for PSR D1133]16 bv inspecting our 10 strong pulses at the ower frequencies., We can study the broadband characteristics for PSR B1133+16 by inspecting our 40 strong pulses at the lower frequencies.1045 The three strongest pulses are shown in Fie. 17.., The three strongest pulses are shown in Fig. \ref{giants}.1046 It turus out that 36 out of the Ü pulses are also mach stronger than the average pulse at a eivenphase at the lower frequencies., It turns out that 36 out of the 40 pulses are also much stronger than the average pulse at a given phase at the lower frequencies.1047 But since the euidssion of these pulses is inostly. concentrated in longitudes of the leaciug component. while the second component is significantly stronger atf the lower frequencies. he mean flux deusitv does not exceed the eiaut-pulse threshold at hose.," But since the emission of these pulses is mostly concentrated in longitudes of the leading component, while the second component is significantly stronger at the lower frequencies, the mean flux density does not exceed the giant-pulse threshold at those."1048 There are [pulses (or 10543). however. which do not appear to be broadband in their strength.," There are 4 pulses (or ), however, which do not appear to be broadband in their strength."1049 Oue of those is shown in the same Figure., One of those is shown in the same Figure.1050 Based ou the very small sample of pulsus kuownu to exhibit giant pulses. it had been suggestedS8 that the occurrence of elant pulses may be related to the magnetic field streneth at the πο cylinder (c.g. Cognard ct al.," Based on the very small sample of pulsars known to exhibit giant pulses, it had been suggested that the occurrence of giant pulses may be related to the magnetic field strength at the light cylinder (e.g. Cognard et al."1051 1996) or to the emission process creating X-ray and eanunuerav endssion (Jolustou Romani 2002)., 1996) or to the emission process creating X-ray and gamma-ray emission (Johnston Romani 2002).1052 While PSR Bl133116 is not amone the list of detected X-rav cluitters. its magnetic field at the light evliuder is five orders of magnitude simaller than that of the Crab mulsar or PSR BL937|21.," While PSR B1133+16 is not among the list of detected X-ray emitters, its magnetic field at the light cylinder is five orders of magnitude smaller than that of the Crab pulsar or PSR B1937+21."1053 Overall. however. if seclus )ossible hat the strong pulses seen for PSR D1133|16 i 1554 MITZ are a population of giaut pulses which ouly uauifest themselves at the higher frequeucies because heir radio spectrüni is significauflv flatter than that of the normal pulsar cussion.," Overall, however, it seems possible that the strong pulses seen for PSR B1133+16 at 4850 MHz are a population of giant pulses which only manifest themselves at the higher frequencies because their radio spectrum is significantly flatter than that of the normal pulsar emission."1054" This is cousisteut with oservatious of single pulses of PSR D1133|16 at an even [umigher frequency of 8150. MITz where the ""giant pulses appear even more frequently (Alaron Lóbluuer. private conumaimuication)."," This is consistent with observations of single pulses of PSR B1133+16 at an even higher frequency of 8450 MHz where the “giant” pulses appear even more frequently (Maron Löhhmer, private communication)."1055 Certainly. a classification as giant pulses cannot be based on the measured flux deusitv alone (see also Johustou Ποια 2002). so that better statistics are needed to determine as to whether a power-law is present in their cnerey distribution to coufinu their nature.," Certainly, a classification as giant pulses cannot be based on the measured flux density alone (see also Johnston Romani 2002), so that better statistics are needed to determine as to whether a power-law is present in their energy distribution to confirm their nature."1056 Analvsing simultaneous multi-frequency observations of PSRs Bo329)51 and D11331]16 we have derived iutrinsic flux. density tine series by correcting for scintillation effects., Analysing simultaneous multi-frequency observations of PSRs B0329+54 and B1133+16 we have derived intrinsic flux density time series by correcting for scintillation effects.1057 The resulting pulse-to-pulse modulation indices show a nuüunununi around 1 Cz and inerease above aud below this frequency., The resulting pulse-to-pulse modulation indices show a minimum around 1 GHz and increase above and below this frequency.1058 This effect may be caused by a loss iu coherence., This effect may be caused by a loss in coherence.1059 Iudecd. neispectiue the correlations between the fiux densities of ~ifereut frequency pairs. it is striking that the correlation OCYOASCN απ the frequency widens.," Indeed, inspecting the correlations between the flux densities of different frequency pairs, it is striking that the correlation decreases as the frequency widens."1060 Snadlar effects ave been observed for the polarization properties in Iouwastereiou et al. (, Similar effects have been observed for the polarization properties in Karastergiou et al. (10612001. 2002) and are clear indications or an intrinsic bandwidth of the enuüssion process or xopagation effects in the maguetosphere.,"2001, 2002) and are clear indications for an intrinsic bandwidth of the emission process or propagation effects in the magnetosphere."1062 We determined the intrinsic flux density spectra or single pulses of PSR BoO329(51 aud derive similar conclusious as obtained from the average profiles., We determined the intrinsic flux density spectra for single pulses of PSR B0329+54 and derive similar conclusions as obtained from the average profiles.1063 This sugecst a relatively stable formation xocess of the average o»ilse profile. with spectral iudices which are flatter for he ceutral part than for the outer componcuts.," This suggest a relatively stable formation process of the average pulse profile, with spectral indices which are flatter for the central part than for the outer components."1064 The, The1065tational processes taking place during the cluster orluation. as well as auv additional energy input. e.g. from a central heat eugimoe. aud these processes are expected to differ significantly frou structure o structure.,"tational processes taking place during the cluster formation, as well as any additional energy input, e.g. from a central heat engine, and these processes are expected to differ significantly from structure to structure."1066 The eas density profile also exhibits a roughly universal behaviour. xahich has allowed observers to fit remarkably simple forms. c.g. a veta-profile (Cavaliere&Fusco-FenuianoSarazin 1986).. to the data.," The gas density profile also exhibits a roughly universal behaviour, which has allowed observers to fit remarkably simple forms, e.g. a beta-profile \citep{1976A&A....49..137C,1986RvMP...58....1S}, to the data."1067 A natural question arises as to whether the xoperties of the hot N-vav cutting σας in galaxy clusters can bepredicted from. first principles starting from the cluster DM distribution alone.s," A natural question arises as to whether the properties of the hot X-ray emitting gas in galaxy clusters can be from first principles starting from the cluster DM distribution alone.,"1068pecifically. we will show how the gas density profile can be obtained cirectly from. the uuderhiug DAL profile. bv combining the equation of livcdyostatic equilibria for the eas and the Jeaus equation for the DAsvuuuetry.," we will show how the gas density profile can be obtained directly from the underlying DM profile, by combining the equation of hydrostatic equilibrium for the gas and the Jeans equation for the DM."1069 Besides the aboveaneutioned relation between the eas and DAL temperatures. our derivation rests upon a very simple connection between the DM velocity anisotropy aud the slope of its density profile. which has receutlv eierged in umunucercal siuulations (Tausen&Moore2006:HansenStadel 2006).," Besides the above-mentioned relation between the gas and DM temperatures, our derivation rests upon a very simple connection between the DM velocity anisotropy and the slope of its density profile, which has recently emerged in numerical simulations \citep{hansenmoore,hansenstadel}."1070.. We thereby demoustrate that the eas deusitv profile is completely deteriiued once the eravitationally donnant DAL deusity profile is even., We thereby demonstrate that the gas density profile is completely determined once the gravitationally dominant DM density profile is given.1071 Since the eas temperature profile is also known. it turus out that the DM cistribution dictates all the eas propertiesιοί.," Since the gas temperature profile is also known, it turns out that the DM distribution dictates all the gas properties."1072 Besides conceptually relevant in itself. this fact allows to predict the ταν huuiuositv xofile of a cluster in terms of its DAL conteut alone(1998).," Besides conceptually relevant in itself, this fact allows to predict the X-ray luminosity profile of a cluster in terms of its DM content alone."1073 So. a new strateev becomes available to coustrain the DM morphology iu galaxy clusters from X-rav observations.," So, a new strategy becomes available to constrain the DM morphology in galaxy clusters from X-ray observations."1074 Moreover. our findings can be eniploved as a practical tool for creating initial conditions for realistic cosmological structures to be used in nunuerical simulations.," Moreover, our findings can be employed as a practical tool for creating initial conditions for realistic cosmological structures to be used in numerical simulations."1075 We start by recalling some basic iformation which will be instrmmental for our analysis., We start by recalling some basic information which will be instrumental for our analysis.1076 We restrict our attention throughout to regular clusters. which are supposed to be spherically svuunetric and relaxed.," We restrict our attention throughout to regular clusters, which are supposed to be spherically symmetric and relaxed."1077" The coudition of hvdrostatic equilibrium for the X-ray cinitting gas can be written as where p4(r) aud Tj(r) ave the gas density aud eniperature profiles. respectively. µ20.61 is the uean inoleculay weight for the intracluster gas. nn, is the proton inass and Mys(r) represcuts he total mass inside radius r."," The condition of hydrostatic equilibrium for the X-ray emitting gas can be written as where $\rho_g (r)$ and $T_g (r)$ are the gas density and temperature profiles, respectively, $\mu \simeq 0.61$ is the mean molecular weight for the intracluster gas, $m_p$ is the proton mass and $M_{\mathrm{tot}}(r)$ represents the total mass inside radius $r$."1078 Two couditious rave to be satisfied in order for Eq. (1)), Two conditions have to be satisfied in order for Eq. \ref{a1}) )1079 o hold., to hold.1080 First. it should be applied to a region cousiderably arecr than the eas mean free path. so that ocal thermodynamic equilibrimia is established.," First, it should be applied to a region considerably larger than the gas mean free path, so that local thermodynamic equilibrium is established."1081 Secoud. the cooling time in that region should be arecr than the age of the cluster. so that no bulk notion occurs.," Second, the cooling time in that region should be larger than the age of the cluster, so that no bulk motion occurs."1082 The latter condition is generally uet outside the central region. where the preseuce of a cooling flow often requires Eq. (1))," The latter condition is generally met outside the central region, where the presence of a cooling flow often requires Eq. \ref{a1}) )"1083 to be replaced by the Euler equation (with the velocity enu plaving a nonnuceligible role)., to be replaced by the Euler equation (with the velocity term playing a nonnegligible role).1084 Because of collisional relaxation. the eas velocity distribution is Isotropic and its temperature can be expressed in terms of the ouc-cimensional velocity dispersion Asstuning complete spherical svnuuetry for the DAL distribution. the two tangential componcuts of the DAL velocity dispersion. denoted by a. are necessarily equal. but they are generally allowed to differ frous the radial commpoucut 62. since DM is supposed to be collisionless.," Because of collisional relaxation, the gas velocity distribution is isotropic and its temperature can be expressed in terms of the one-dimensional velocity dispersion ${\sigma_g^2}$ as Assuming complete spherical symmetry for the DM distribution, the two tangential components of the DM velocity dispersion, denoted by ${\sigma}_t^2$, are necessarily equal, but they are generally allowed to differ from the radial component ${\sigma}_r^2$, since DM is supposed to be collisionless."1085" It is usual to quantity the DM velocity anisotropy by aud we fudit conveuieut to iutroduce the mean DM one-dimensional. velocity. dispersion. oj,, as", It is usual to quantify the DM velocity anisotropy by and we find it convenient to introduce the mean DM one-dimensional velocity dispersion $\sigma_{\mathrm{DM}}^2$ as1086approximations are needed to solve the “Fermjon sign problemi in cases where exchange interaction becomes essential.,"approximations are needed to solve the ""Fermion sign problem"" in cases where exchange interaction becomes essential."1087 The temperature dependent mixed state description of the I. jon. the density dependent equilibriua dissociation recombination balance and the euergeties js been evaluated for the first tine.," The temperature dependent mixed state description of the $_3^+$ ion, the density dependent equilibrium dissociation recombination balance and the energetics has been evaluated for the first time."1088 With the rising cluperature the rovibrational excitations contribute o the cnerectics. as expected. whereas the electronic wart rendus iu its eround state in the spirit of the BornOppeculeiner approximation.," With the rising temperature the rovibrational excitations contribute to the energetics, as expected, whereas the electronic part remains in its ground state in the spirit of the Born–Oppenheimer approximation."1089 At about 1000 IX he fragments of the iuolecule. ID|II!. IT]LIT. and 2]E LIT. start coutributiug.," At about $4000$ K the fragments of the molecule, $_2+$ $^+$, $_2^++$ H and $2$ $+$ $^+$, start contributing."1090 Therefore. IL] ion becomes ess dominant. and eventually negligible in high euough T," Therefore, $_3^+$ ion becomes less dominant, and eventually negligible in high enough $T$."1091 We have shown how the partial decohereuce iu the mixed state can be used for interpretation of the fragment composition of the equilibrium reaction., We have shown how the partial decoherence in the mixed state can be used for interpretation of the fragment composition of the equilibrium reaction.1092 Furthermore. we have evaluated explicitly the related partition fiction. free cnerey. entropy. and heat capacity. all as functions of temperature.," Furthermore, we have evaluated explicitly the related partition function, free energy, entropy and heat capacity, all as functions of temperature."1093 An accurate analytical fictional foris are given for the temperatures below dissociation., An accurate analytical functional forms are given for the temperatures below dissociation.1094 We consider all these as major improvements to the eurlier published studies. where dissociation has not Όσοι considered.," We consider all these as major improvements to the earlier published studies, where dissociation has not been considered."1095 For financial support we thank the Academy of Fiulaud. aud for computational resources the facilities of Fiuuish ΤΕ Center for Science (CSC) aud Material Sciences National Grid Iufrastructure (ALerid. akaatti).," For financial support we thank the Academy of Finland, and for computational resources the facilities of Finnish IT Center for Science (CSC) and Material Sciences National Grid Infrastructure (M-grid, akaatti)."1096presented bv Deutsch (1956) who found circuustellar lines iu the spectrum of the visual companion to the «πηρασπααν,presented by Deutsch (1956) who found circumstellar lines in the spectrum of the visual companion to the supergiant.1097 Wevinanus iuferenees about the shell size were compatible with the direct lower Iluuit to the outer shell radius from Deutscl’s observation., Weymann's inferences about the shell size were compatible with the direct lower limit to the outer shell radius from Deutsch's observation.1098 Observations of 21 cmi cutission from Detelgeuse have shown that gas extends to about 1 minute of arc from the star (Bowers Knapp 19857)., Observations of 21 cm emission from Betelgeuse have shown that gas extends to about 1 minute of arc from the star (Bowers Knapp 1987).1099 Betelgeuse’s circtuustellar shell was resolved spatially by Bernat Lambert (1975) who detected fluorescent clnission iu the IK I resonance line at augular distances of 2 - [| arcsec frou the star., Betelgeuse's circumstellar shell was resolved spatially by Bernat Lambert (1975) who detected fluorescent emission in the K I resonance line at angular distances of 2 - 4 arcsec from the star.1100 Subsequent observations extended the detections of the Cluission to abou GU arcsec from the star and provided evidence for some clepartures from spherical svuuuetiv aud for structures 1 the shell (Bernat Laihbert 1976: Bernat et al., Subsequent observations extended the detections of the emission to about 60 arcsec from the star and provided evidence for some departures from spherical symmetry and for structures in the shell (Bernat Lambert 1976; Bernat et al.1101 1978: IToueveutt et al., 1978; Honeycutt et al.1102 1980: AMlauron et al., 1980; Mauron et al.1103 1981: Miinwou 1990), 1984; Mauron 1990).1104 Enuission iu the Na D lines was detected by Mauroji Οπως (1990). Maurou. (1990). and AMlauron Coulain (1995).," Emission in the Na D lines was detected by Mauron Querci (1990), Mauron (1990), and Mauron Guilain (1995)."1105 The most remote (from the star) manifestaion of the ciremustellar shell is the parsec-sized bow shock reported by Noricga-Crespo ct al. (, The most remote (from the star) manifestation of the circumstellar shell is the parsec-sized bow shock reported by Noriega-Crespo et al. (11061997) from IRAS nuages at 60 and LOO pan. Exiission detected from an are with a mean radius of 6 minutes of are is attributed to material confined by the interstellar medis ran pressure.,1997) from IRAS images at 60 and 100 $\mu$ m. Emission detected from an arc with a mean radius of 6 minutes of arc is attributed to material confined by the interstellar medium's ram pressure.1107" Iu this paper. we present aud analyse a new series of observations of eecnmission obtained with a long slit at a resolving power of about 110.000 or 2.6 kins 1 with the slit placed at various Hupact parameters frou, the star aud position angles around the stax. and at several epochs."," In this paper, we present and analyse a new series of observations of emission obtained with a long slit at a resolving power of about 110,000 or 2.6 km $^{-1}$ with the slit placed at various impact parameters from the star and position angles around the star, and at several epochs."1108 Our combination o spectral aud spatial resolution was not achieved by carlicr studies of the cnussion., Our combination of spectral and spatial resolution was not achieved by earlier studies of the emission.1109 We derive afresh the radial dependence of the emission and rediscuss the principal conclusion - the mass loss rate - that has been previously drawn from the radial dependence., We derive afresh the radial dependence of the emission and rediscuss the principal conclusion - the mass loss rate - that has been previously drawn from the radial dependence.1110 A novel result of our observations --s the discovery of chimps iu the cussion and. hence. oei the neutral eas in the shell aud. wind.," A novel result of our observations is the discovery of clumps in the emission and, hence, in the neutral gas in the shell and wind."1111 Some of these oethomogenecitics are unresolved both spatially iu our secing-limited observations and also kinematically at the limit of about 2 kms. |., Some of these inhomogeneities are unresolved both spatially in our seeing-limited observations and also kinematically at the limit of about 2 km $^{-1}$.1112 Detelgeuse was observed in 1991 February. March. aud November with the W.J. MeDouald Observatorvs IEbulau J. Suüth 2.710 reflector aud its coudé spectrograph.," Betelgeuse was observed in 1994 February, March, and November with the W.J. McDonald Observatory's Harlan J. Smith 2.7m reflector and its coudé spectrograph."1113 The latter was used with au echelle erating and an interference filter to separate the order providing the Hine., The latter was used with an echelle grating and an interference filter to separate the order providing the line.1114 The slit length was varied with a majority of the observations taken with a length of 10 to 100 arcsec., The slit length was varied with a majority of the observations taken with a length of 10 to 100 arcsec.1115 The slit width was almost always set at 0.55 arcsec., The slit width was almost always set at 0.55 arcsec.1116 Spectra were recorded bv a Tektronix 512 X 512 CCD over a bandpass of about17À., Spectra were recorded by a Tektronix 512 X 512 CCD over a bandpass of about.1117. Standard reduction procedures were applied to correct for the bias and pixcl-to-pixel variations., Standard reduction procedures were applied to correct for the bias and pixel-to-pixel variations.1118 The waveleneth scale was set using Thi-Àr lamp spectra., The wavelength scale was set using Th-Ar lamp spectra.1119 During cach observing session. several ou-star spectra were recorded with a very short slit leugth.," During each observing session, several on-star spectra were recorded with a very short slit length."1120 Offstar spectra were then collected with longer slit leugths., Off-star spectra were then collected with longer slit lengths.1121 Iu eecneral the star was not included iu the slit: inclusion limited the exposure time bevoud which the starlight Neurated a portion of the CCD chip.," In general, the star was not included in the slit; inclusion limited the exposure time beyond which the starlight saturated a portion of the CCD chip."1122" Since the fell at the coudé focus rotaes, exposure times were kept short to ninindze sniearnug of enüssion inlomogeneities: typically. exposure tines were 10 minutes or less."," Since the field at the coudé focus rotates, exposure times were kept short to minimize smearing of emission inhomogeneities; typically, exposure times were 10 minutes or less."1123 For capturing chussion far frou the star. a few exposures as long as 30 münuutes were made.," For capturing emission far from the star, a few exposures as long as 30 minutes were made."1124 Relative spectroplotometiry was attempted: observations were terminated im partly cloudy skies., Relative spectrophotometry was attempted; observations were terminated in partly cloudy skies.1125 The stellar flux was measured) by openine he slit ο. its mnaxinun width when all the visible image was contained within the slit., The stellar flux was measured by opening the slit to its maximum width when all the visible image was contained within the slit.1126 A euider maintained the stars )osition relative to the slit to an accuracy of about 0.5 aresec., A guider maintained the star's position relative to the slit to an accuracy of about 0.5 arcsec.1127 Our reduction tecliniques were adapted directly fro hose described earlier (Plez Lambert 1991)., Our reduction techniques were adapted directly from those described earlier (Plez Lambert 1994).1128 In short. he spectrum at a eiven location off the star is presumed o be a composite of emission from the gas down the line of sieht through the shell at that poiut aud starlight scater in the Earth's atmosphere aud the telescope.," In short, the spectrum at a given location off the star is presumed to be a composite of emission from the gas down the line of sight through the shell at that point and starlight scatter in the Earth's atmosphere and the telescope."1129 We assume hat the latter coutributiou may be represented by a scaled version of aui on-star spectrum., We assume that the latter contribution may be represented by a scaled version of an on-star spectrum.1130 Fie., Fig.1131 1 provides an example of a raw offstar spectrum ac the spectral mage after subtraction of the photospheric scaered ight., \ref{fig1} provides an example of a raw off-star spectrum and the spectral image after subtraction of the photospheric scattered light.1132" For positions of the slit Lear the star (a few arcsec), the plotospheric scatοχος contribution larecly dominates over the light scattered by the shell."," For positions of the slit near the star (a few arcsec), the photospheric scattered contribution largely dominates over the light scattered by the shell."1133" This makes the subtraction difficult. as the inperfectlv removed background spectruni shows in Fig.o 1,"," This makes the subtraction difficult, as the imperfectly removed background spectrum shows in Fig. \ref{fig1}."1134 The terrestrial Os lines appear weakly iu ciission in the image after subtraction., The terrestrial $_2$ lines appear weakly in emission in the image after subtraction.1135 This is due to the difference iu iluuination across the slit when the star is ou one side of the slit., This is due to the difference in illumination across the slit when the star is on one side of the slit.1136 For spectra taken very near the star. we could improve subtraction of the photospheric component by shitting the spectra by less than half a pixel iu the dispersion direction.," For spectra taken very near the star, we could improve subtraction of the photospheric component by shifting the spectra by less than half a pixel in the dispersion direction."1137 Sheht changes in airimass between the ou-star aud οΤατ exposures iav also affect the subtraction of the O» lines., Slight changes in airmass between the on-star and off-star exposures may also affect the subtraction of the $_2$ lines.1138 The spectruii after subtraction of the scattered ight is a velocitv-position map of the emission., The spectrum after subtraction of the scattered light is a velocity-position map of the emission.1139 Rows of pixels xwallel to the dispersion give the spectrum at different outs im the shell., Rows of pixels parallel to the dispersion give the spectrum at different points in the shell.1140 The radial distance of a poiut from Deteleeuse may be calculated from the inmipact parameter (3 avescee in this case) aud the distance in seconds of arc roni the point on the slit of closest approach to the star., The radial distance of a point from Betelgeuse may be calculated from the impact parameter (3 arcsec in this case) and the distance in seconds of arc from the point on the slit of closest approach to the star.1141 The orientation of the slit was derived from the hour anele of the observation., The orientation of the slit was derived from the hour angle of the observation.1142 A thin uniform shell of outwardly Howiue and fluorescing atomis appears as an cllipse in such a velocity-position map., A thin uniform shell of outwardly flowing and fluorescing atoms appears as an ellipse in such a velocity-position map.1143 Such an elliptical looking, Such an elliptical looking1144two stars.,two stars.1145 Olumic dissipation of these currents in the flux tubes’ footpoiuts on the maguectic star gives rise to the N-ravs., Ohmic dissipation of these currents in the flux tubes' footpoints on the magnetic star gives rise to the X-rays.1146 In this model the 5.[-13i1uute. period is the orbital period. but the two stars are detached.," In this model the 5.4-minute period is the orbital period, but the two stars are detached."1147 Since the binary loses angular momentum due to eravitational-wave radiation. it is expected to evolve towards shorter orbital periods. as observed.," Since the binary loses angular momentum due to gravitational-wave radiation, it is expected to evolve towards shorter orbital periods, as observed."1148 The main problem is the offset between the optical aud N-ray flux. which requires the footpoiuts on the maenetic stars surface to be almost 90 degrees ahead of the orbiting non-magnetic star iu both svstems (Barrosetal.2005.2007).," The main problem is the offset between the optical and X-ray flux, which requires the footpoints on the magnetic star's surface to be almost 90 degrees ahead of the orbiting non-magnetic star in both systems \citep{Bar05,Bar07}."1149. The third. AM. ΟΝα) model. has a Roche-lobe filling white dwarf losing mass to a more massive white dwarf (Alarsh&Steceghs2002:Ramsayetal.2002b.. or Cropperetal.1998. for a magnetic version).," The third, `AM CVn' model, has a Roche-lobe filling white dwarf losing mass to a more massive white dwarf \citealt{Mar02,Ram02b}, or \citealt{Cro98} for a magnetic version)."1150 The orbital period is 5. Laninutes in the case ofCnc.. aud the accretion stream hits the aceretor directly without forming a disk.," The orbital period is 5.4 minutes in the case of, and the accretion stream hits the accretor directly without forming a disk."1151 The phase offset between optical and N-rav flux is naturally accounted for. since the accretion stream will deflect and lit the accretor off-axis.," The phase offset between optical and X-ray flux is naturally accounted for, since the accretion stream will deflect and hit the accretor off-axis."1152 The absence of longer periods is also expected. but the observed decrease of the periods in aand VLO? Vul is cousidered problematic for this model (Strolunaver2002.2001). although solutions to this problem have been put forward (Delove&Taam2006:D'Antonaetal. 2006:: Roclofs Delove in preparation).," The absence of longer periods is also expected, but the observed decrease of the periods in and V407 Vul is considered problematic for this model \citep{Str02,Str04}, although solutions to this problem have been put forward \citealt{Del06,DAn06}; Roelofs Deloye in preparation)."1153" This Letter describes our search for kincmatic evidence, using pliase-resolved spectroscopy ofCuc.. for or against these different models."," This Letter describes our search for kinematic evidence, using phase-resolved spectroscopy of, for or against these different models."1154 Section refsec:obs describes our observations auddata reduction., Section \\ref{sec:obs} describes our observations and data reduction.1155 We preseut our results in refseciesults and discuss their implications iu refseciconclusion.., We present our results in \\ref{sec:results} and discuss their implications in \\ref{sec:conclusion}.1156 After wuusuccessful attempts in 2005. 2006 2007 to obtain plase-resolved spectroscopy ofCie.. we finally inet acceptable weather conditions at νους ou January 25 26 and March 30 of 2009.," After unsuccessful attempts in 2005, 2006 2007 to obtain phase-resolved spectroscopy of, we finally met acceptable weather conditions at Keck-I on January 25 26 and March 30 of 2009."1157" We took spectra using a 1.5"" slit aud the 600 erooves/1uau eris ou the blue side of LRIS. the Low-Resolution huaeiug Spectrograph (Okeetal.1995).. for an effective spectral resolution of about 300 Xiun/s. Di order to resolve the 321.5-3ecoud period in Ciuc... we linited the exposure times to G0 seconds;"," We took spectra using a $''$ slit and the 600 grooves/mm grism on the blue side of LRIS, the Low-Resolution Imaging Spectrograph \citep{Oke95}, for an effective spectral resolution of about 300 km/s. In order to resolve the 321.5-second period in , we limited the exposure times to 60 seconds."1158" The blue CCD. with a plate scale of 0.135"" /pixel. was binuecd Lb«&{ pixels to reduce the read-out noise while still sampling resolution elements with approximately 2 (binned) pixels."," The blue CCD, with a plate scale of $''$ /pixel, was binned $4\times 4$ pixels to reduce the read-out noise while still sampling resolution elements with approximately 2 (binned) pixels."1159 The biuniung firthermore reduced the dead-time between exposures| to 27 seconds., The binning furthermore reduced the dead-time between exposures to 27 seconds.1160 Data reduction was done using staudard tasks., Data reduction was done using standard tasks.1161 To optimally extract the very faint individual spectra hhas optical maenitudes C~19.6.Bx20.7: Israelet. 2002)). reference aperture profiles were created from averaged blocks of typically —20 successive frames to increase the photon statistics and remove cosnic ravs from the profile.," To optimally extract the very faint individual spectra has optical magnitudes $U\approx 19.6, B\approx 20.7$; \citealt{Isr02}) ), reference aperture profiles were created from averaged blocks of typically $\sim$ 30 successive frames to increase the photon statistics and remove cosmic rays from the profile."1162 Spectra were grouped together based ou the requirement that spatial ceuter and profile (sccine) ou the CCD be (near-jicdentical., Spectra were grouped together based on the requirement that spatial center and profile (seeing) on the CCD be (near-)identical.1163 Wavelength calibration was achieved using TeCCdZZuNNcAAr are exposures., Wavelength calibration was achieved using Ar arc exposures.1164 A fit of 30 arc lines left irrootanean-square residuals., A fit of 30 arc lines left root-mean-square residuals.1165 No detectable arc drift occurred during the nieht. aud the dift between the two nights was well below oue pixcl.," No detectable arc drift occurred during the night, and the drift between the two nights was well below one pixel."1166 We thus applied a conumion dispersion solution to all spectra. only applviug a σα. shift to the first nights spectra to tie the strong 1] 5577 sky cinissiou lines together.," We thus applied a common dispersion solution to all spectra, only applying a small shift to the first night's spectra to tie the strong ] 5577 sky emission lines together."1167 The relative wavelength[O calibration of all spectra is estimated to be better than0., The relative wavelength calibration of all spectra is estimated to be better than.11681À.. The time stamp and velocity scale of cach spectrum were transformed to the solar systems barvceeuter., The time stamp and velocity scale of each spectrum were transformed to the solar system's barycenter.1169 Time stamps were checked several times durus our run using Universal Time clocks available ouline aud estimated. to be correct. to one second. (or better)., Time stamps were checked several times during our run using Universal Time clocks available online and estimated to be correct to one second (or better).1170 The spectrophotometric standard star Feige 231 was used for calibrating the iustruneutal response: variable transparency and ποστς prechided an absolute fiux calibration., The spectrophotometric standard star Feige 34 was used for calibrating the instrumental response; variable transparency and seeing precluded an absolute flux calibration.1171 Table 1. sununuarizes our observiug log., Table \ref{table:log} summarizes our observing log.1172 The average Neck spectrum of lis shown in reffie:spectimm.., The average Keck spectrum of is shown in \\ref{fig:spectrum}.1173 As previously reported (Israeletal.2002) the spectrmi is clominated bv ionized reli cussion lines., As previously reported \citep{Isr02} the spectrum is dominated by ionized helium emission lines.1174 Their full-idth at half£aunaxinma (FWIDBND of ~2500kkinss | is well resolved bv the + resolution of our spectra., Their full-width at half-maximum (FWHM) of $\sim$ $^{-1}$ is well resolved by the $^{-1}$ resolution of our spectra.1175 Lines of neutral icliuni are also present but much weaker., Lines of neutral helium are also present but much weaker.1176 Our ligher-quality spectruu confiiis earlier fudiues that the eveu-erm transitions of thei1 Pickering series are stronger han the odd-terui ones (see refüe:spectrumn)). which was iuterpreted by Nortouctal.(2001). aud. Boiuschetal.2007). as evidence for he presence of hwdrogen.," Our higher-quality spectrum confirms earlier findings that the even-term transitions of the Pickering series are stronger than the odd-term ones (see \\ref{fig:spectrum}) ), which was interpreted by \citet{Nor04} and \citet{Rei07} as evidence for the presence of hydrogen."1177 Iu particular the absence ofTet 1200 in our spectrum is striking. andmost casily explained if oue assumes that there is indeed hvdrogenu in (or simplicity. we will refer toHei Pickering lines in the remainder of this paper).," In particular the absence of 4200 in our spectrum is striking, andmost easily explained if one assumes that there is indeed hydrogen in (for simplicity, we will refer to Pickering lines in the remainder of this paper)."1178 Tn addition to the 1610 line in the Bowen bleud identified by Isracletal.(2002).. the strong L379 line appears to be present to the red of £338.," In addition to the 4640 line in the Bowen blend identified by \citet{Isr02}, the strong 4379 line appears to be present to the red of 4338."1179 Appareut enission lines near aand ecould be due toNi. while a faint feature to the red of τά may be 511.," Apparent emission lines near and could be due to, while a faint feature to the red of 4471 may be 4514."1180 The preseuce of several uitrogeu lines appears to be at odds with the suggestion (based on fits to N-rav spectra) that nuuav have an unusual chemical abundance pattern aud be uuderabundaut in nitrogen in particular (Strolumayer 2008)., The presence of several nitrogen lines appears to be at odds with the suggestion (based on fits to X-ray spectra) that may have an unusual chemical abundance pattern and be underabundant in nitrogen in particular \citep{Str08}.1181 We note that uitrogen-rich matter would be expected due to the CNO-cveleif lis an ultracompact binary. unless the imaterial experienced helm buruiug (Yuugelson 2008).. aud that nitrogen indeed appears to be abundant in most AM CVin stars (κου Roclotsetal. 2009)).," We note that nitrogen-rich matter would be expected due to the CNO-cycleif is an ultracompact binary, unless the material experienced helium burning \citep{Yun08}, , and that nitrogen indeed appears to be abundant in most AM CVn stars (see \citealt{Roe09}) )."1182 When we phase-fold our spectra using the ephemeris of Barrosetal. (2007).. we observe continuuniflux," When we phase-fold our spectra using the ephemeris of \citet{Bar07}, , we observe continuumflux"1183»ointings using the phased array. one at the assumed pulsar »osition and the other four olfset from it by hall of WILL in the north. south. cast and west directions.,"pointings using the phased array, one at the assumed pulsar position and the other four offset from it by half of FWHM in the north, south, east and west directions."1184 Phe pulsar was usually detected in at least two of the grid. pointings., The pulsar was usually detected in at least two of the grid pointings.1185 The »osition of the grid pointing was weighted by the measured S/N of detection to arrive at à more refined. position., The position of the grid pointing was weighted by the measured S/N of detection to arrive at a more refined position.1186 The ocalization accuracy was improved by a grid. using first the survey setup with ENIM of 40° and subsequent grids with each of the three selected. phased array configurations. with EFWHIAIs of 20°. 105 and 5°.," The localization accuracy was improved by a grid using first the survey setup with FWHM of 40' and subsequent grids with each of the three selected phased array configurations, with FWHMs of 20', 10' and 5'."1187 The position of the pulsar was refined. alter every. iteration., The position of the pulsar was refined after every iteration.1188 An experiment with known »ulsars of varving flux densities was carried out to test this oocedure., An experiment with known pulsars of varying flux densities was carried out to test this procedure.1189 I was observed that this procedure vielded. the inal positions with up to 3 accuracy in about 7 hours of observations for a weak pulsar (Flux. density at 626 MIIZz of about 2 mJv)., It was observed that this procedure yielded the final positions with up to 3' accuracy in about 7 hours of observations for a weak pulsar (Flux density at 626 MHz of about 2 mJy).1190 The required. observing time for this oocecdure is less for a pulsar with higher flux density., The required observing time for this procedure is less for a pulsar with higher flux density.1191 PSR. 100560|6320 is verv close to the beam centre in he discovery observations. so its position was refined using a grid pointing procedure similar to that described above using the Lovell telescope at Jodrell Bank Observatory at a ligher observing frequeney of LA Cllz.," PSR J0026+6320 is very close to the beam centre in the discovery observations, so its position was refined using a grid pointing procedure similar to that described above using the Lovell telescope at Jodrell Bank Observatory at a higher observing frequency of 1.4 GHz."1192 The new localization oocedure described here was very useful for the other two oulsars. particularly for PSR. J2217|5733. as these pulsars were discovered at the edge of the beam.," The new localization procedure described here was very useful for the other two pulsars, particularly for PSR J2217+5733, as these pulsars were discovered at the edge of the beam."1193 Once a position accurate to approximately 3° was obtained using the procedure. outlined above. the new pulsar was observed. using 762m. Lovell telescope at Jodrell Bank Observatory on multiple epochs to obtain a timing solution.," Once a position accurate to approximately 3' was obtained using the procedure outlined above, the new pulsar was observed using 76-m Lovell telescope at Jodrell Bank Observatory on multiple epochs to obtain a timing solution."1194 The Lovell telescope is equipped with a cual channel Crvogenic receiver with a receiver temperature of about 35 Ix zt. 1400. MlIzZz., The Lovell telescope is equipped with a dual channel cryogenic receiver with a receiver temperature of about 35 K at 1400 MHz.1195 This receiver is sensitive to two hands of circular polarizations., This receiver is sensitive to two hands of circular polarizations.1196 The radio [frequency signals are down-converted to an intermediate frequency. at. the focus of the telescope ancl brought to the receiver room. where hese are further clown-convertec before being fed to the oulsar hardware.," The radio frequency signals are down-converted to an intermediate frequency at the focus of the telescope and brought to the receiver room, where these are further down-converted before being fed to the pulsar hardware."1197 Ehe pulsars were observed. with a 32-Παπο clual hardware filterbank across a 32 MllIz bandpass entered around. 1400 MIIz., The pulsars were observed with a 32-channel dual hardware filterbank across a 32 MHz bandpass centered around 1400 MHz.1198 The total intensity data were yen cleclispersed in a hardware cleclispersion unit and folded synchronously at the. nominal topocentric period. of the »ulsar for subintegrations of between 1 to 2 minutes., The total intensity data were then dedispersed in a hardware dedispersion unit and folded synchronously at the nominal topocentric period of the pulsar for subintegrations of between 1 to 2 minutes.1199 The olded data for typically 6 to 10 subintegrations were written ο cisk for every epoch and observations were usually made at approximately regular intervals of between 10 to 15 davs., The folded data for typically 6 to 10 subintegrations were written to disk for every epoch and observations were usually made at approximately regular intervals of between 10 to 15 days.1200 A time stamp derived from a hydrogen. maser clock was recorded: with each folded: profile., A time stamp derived from a hydrogen maser clock was recorded with each folded profile.

Showing the first 1,200 of 10420 lines. Download the file for the rest.