ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Open field mes can. and likely do. trausport some angular momenti frou he star and/or disk (via winds).," Open field lines can, and likely do, transport some angular momentum from the star and/or disk (via winds)."3 However. no augular uonmentun is exchanged between the star aud disk aloug open field lines.," However, no angular momentum is exchanged between the star and disk along open field lines."4 Tn the preseut work. we consider oulv he torques arising in the star-disk interaction and thus weelect auv torques that may arise from open field lues.," In the present work, we consider only the torques arising in the star-disk interaction and thus neglect any torques that may arise from open field lines."5 To take iuto account the opening of the field. the \IPO5 ormulation inclides a maxiuimn for the absolute value of the twist. σον," To take into account the opening of the field, the MP05 formulation includes a maximum for the absolute value of the twist, $\gamma_c$."6 Iu the regions of the disk where the twist is greater than this critical value. the mocel assiues that he magnetic ficld no longer connects the star to the disk.," In the regions of the disk where the twist is greater than this critical value, the model assumes that the magnetic field no longer connects the star to the disk."7 The torque on the star that would have arisen from those nagnuetic field lines is mstead taken to be zero., The torque on the star that would have arisen from those magnetic field lines is instead taken to be zero.8" MP05 showed that a value of 5,=ox represcuts the classical asstuuption of a field that remains connected at all radi in the disk.", MP05 showed that a value of $\gamma_c = \infty$ represents the classical assumption of a field that remains connected at all radii in the disk.9 To show the effect of the opening of the magnetic field iu more realistic systems. a value of +.=1 is appropriate (Uzdeuskyetal.2002).," To show the effect of the opening of the magnetic field in more realistic systems, a value of $\gamma_c = 1$ is appropriate \citep{uzdensky3ea02}."10. Furthermore. \IP05 showed that there exists a mode change in the magnetic connection between the star aud disk. at a threshold value of the spin rate.," Furthermore, MP05 showed that there exists a mode change in the magnetic connection between the star and disk, at a threshold value of the spin rate."11 Specifically. if where then the stellar 1nagnetic field ouly connects to a simall region near the inner edee of the disk.," Specifically, if where then the stellar magnetic field only connects to a small region near the inner edge of the disk."12" In this case. which they call ""State 1° no magnetic field connects"," In this case, which they call “State 1,” no magnetic field connects"13seven galaxies Wwath redshift coutrold «i>=0.09001+0.0007 aud velocity dispersion of 53y+200Ians ,"seven galaxies with redshift centroid $<z>=0.0901 \pm 0.0007$ and velocity dispersion of $530 \pm 20014\kms $."15hupact parameter range is 38.5. 117.58| kpc(h=Hifi100lans!Mpο1). gy—=yn 0.5).," Impact parameter range is $38.5$ – $117.8 \ h^{-1}$ kpc $h~=~H_0/(100 \kms \mbox{Mpc}^{-1})$, $q_0=0.5$ )."16 The fact tha we observe a nuuiber of resolved lines shows that the absorptious arise in discrete clouds of cold οἱs dustead of arising iu a large sinele diffuse eas coniponent., The fact that we observe a number of resolved lines shows that the absorptions arise in discrete clouds of cold gas instead of arising in a large single diffuse gas component.17 The comparison between the cross-correlation function (CF) between the eaaxies ancl absor)ers in the data with the CF correspondiug tfo a raloni case shows that tlfare is a non-randoni connection between the absorbers iud t1e Galaxies iu our data., The comparison between the cross-correlation function (CF) between the galaxies and absorbers in the data with the CF corresponding to a random case shows that there is a non-random connection between the absorbers and the galaxies in our data.18 If we assume that galaxies within the real grou yor cliister are distribited in velocity space according to a gaussiau cüstribution. ao1e-to-one match )otwoeen the absorbers aud the galaxies camuot be established with the preseut cata.," If we assume that galaxies within the real group or cluster are distributed in velocity space according to a gaussian distribution, a one-to-one match between the absorbers and the galaxies cannot be established with the present data."19 From these results we cau conclude the ollowiue: (1) The detection of groups of adavsortion lines implies that some oft1e OODSOTvers do actually cluster. (, From these results we can conclude the following: (1) The detection of groups of absorption lines implies that some of the absorbers do actually cluster. (202) The velocity spanned by the OODSOTion lines arising m eroups or custers is consistent with the velocity dispersion of tlic! corresponcdiuo oroup or cluster of ealaxies.,2) The velocity spanned by the absorption lines arising in groups or clusters is consistent with the velocity dispersion of the corresponding group or cluster of galaxies.21 This imiplics that the OODSOTvers arising in those clusters occipy the same region iu space than the eaaxies themsclves. (, This implies that the absorbers arising in those clusters occupy the same region in space than the galaxies themselves. (223) There is no strong preference for παοςers to avoid overdense cuvirouments. (,3) There is no strong preference for absorbers to avoid overdense environments. (231) adavsortion eaused by galaxy groups or clusters arises m discrete clouds of COd gas. rather than iu some hypothetical smoothly distributed cold phase iu the intracluster medi.,"4) absorption caused by galaxy groups or clusters arises in discrete clouds of cold gas, rather than in some hypothetical smoothly distributed cold phase in the intracluster medium."24AN-stales of positive parity. which are forbidden in this reaction.,"$\Delta N$ -states of positive parity, which are forbidden in this reaction."25 In order to exclude these states one needs to consider the A-isohar ecxitation explicitely., In order to exclude these states one needs to consider the $\Delta$ -isobar ecxitation explicitely.26 Another simplest process which allows to probe fundamental properties of NN svstem is photoabsorption on (wo nucleon svstems., Another simplest process which allows to probe fundamental properties of NN system is photoabsorption on two nucleon systems.27 The deuteron photocisinlegration reaction pn is wiclely used as a testing eround for different theoretical models of the NN-interaction. however. much less is known on the photoclisintegration of the diproton. 5{pp}.—pp. or the inverse process of the photoproduction pp—{pp}.>.," The deuteron photodisintegration reaction $\gamma d\to pn$ is widely used as a testing ground for different theoretical models of the NN-interaction, however, much less is known on the photodisintegration of the diproton, $\gamma \{pp\}_s\to pp$, or the inverse process of the photoproduction $pp\to \{pp\}_s\gamma$."28 Whereas in the photodisinegration of the deuteron the M1 magnetic dipole transition dominates al several hundred MeV. through ihe excitation of the A(1232) isobar. in the reaction with the 4S) diproton M-odd multipoles are forbidden due to angular momentum and. parity conservation.," Whereas in the photodisinegration of the deuteron the M1 magnetic dipole transition dominates at several hundred MeV through the excitation of the $\Delta(1232)$ isobar, in the reaction with the $^1S_0$ diproton M-odd multipoles are forbidden due to angular momentum and parity conservation."29 Therefore. (here is no direct. contribution of the intermediate S-wave A.V states in the reaction pp—>{pp}s.," Therefore, there is no direct contribution of the intermediate S-wave $\Delta N$ states in the reaction $pp\to \gamma \{pp\}_s$."30" Non-direct excitation of the ""59 AN state is possible via the E2 transition [0].. but. this contribution is expected to be less important than the \li-trausition."," Non-direct excitation of the $^5S_2$ $\Delta N$ state is possible via the E2 transition \cite{WNA}, but this contribution is expected to be less important than the M1-transition."31 The OPE model of the reaction pp—(ppl; allows to account for the A contributions via the subprocess zp— p., The OPE model of the reaction $pp\to \{pp\}_s\gamma$ allows to account for the $\Delta$ contributions via the subprocess $\pi^0 p\to p\gamma$ .32 The corresponding OPE diagram is similar to those in Fig. 2..," The corresponding OPE diagram is similar to those in Fig. \ref{fig1},"33" but with the subproscess z""p—po in the down vertex.", but with the subproscess $\pi^0 p\to p\gamma$ in the down vertex.34 The result of the OPE calewlations are shown in Fig.4.., The result of the OPE calculations are shown in \ref{figgamma}.35 One can see that Chis model explaines the observed in Ref., One can see that this model explaines the observed in Ref.36 [4) rise of the cross section almost equantativelv.," \cite{komar08}37 rise of the cross section almost quantatively."38 The second bump at 1.6 GeV is caused by the energy dependence of the πρ—ps cross section [11]. andrelated to excitation of more heavy nucleon isobars., The second bump at 1.6 GeV is caused by the energy dependence of the $\pi^0 p\to p\gamma$ cross section \cite{arndt} andrelated to excitation of more heavy nucleon isobars.39Panels (a) and (c) of Fig.,Panels (a) and (c) of Fig.40" 5 show a portion of the observed penumbra as seen in the continuum intensity (att=449—300 mA)). before and after straylight removal, respectively."," 5 show a portion of the observed penumbra as seen in the continuum intensity (at $\lambda=\lambda_0 - 300$ ), before and after straylight removal, respectively."41 The corresponding maps of Doppler velocity are plotted in panels (b) and (d)., The corresponding maps of Doppler velocity are plotted in panels (b) and (d).42" Obviously both bluc- and redshifts are present in the penumbra, in particular after straylight removal."," Obviously both blue- and redshifts are present in the penumbra, in particular after straylight removal."43 A striking feature in Fig., A striking feature in Fig.44" 5b is the localized patches of strong blueshift, up to 3.3 ss' coinciding with the bright heads of penumbral filaments."," 5b is the localized patches of strong blueshift, up to 3.3 $^{-1}$, coinciding with the bright heads of penumbral filaments."45 These bright heads are surrounded by lanes of gas nearly at rest or slightly redshifted., These bright heads are surrounded by lanes of gas nearly at rest or slightly redshifted.46" In panel (b) significant redshifts are visible only at two locations, once at the side of a filament reaching into the umbra (at x=6"" and y=1"") and once, rather weakly around the bright and strongly blueshifted head of a filament (at x=6.5"" and v=4.5"")."," In panel (b) significant redshifts are visible only at two locations, once at the side of a filament reaching into the umbra (at $x = 6\varcsec$ and $y = 1\varcsec$ ) and once, rather weakly around the bright and strongly blueshifted head of a filament (at $x = 6.5\varcsec$ and $y = 4.5\varcsec$ )."47" A number of new redshifted patches are found in the Doppler map after straylight removal (panel d), appearing dominantly at locations previously seemingly at rest in panel (b)."," A number of new redshifted patches are found in the Doppler map after straylight removal (panel d), appearing dominantly at locations previously seemingly at rest in panel (b)."48" We have grayed out areas where intensity is below 0.6 times that of the quiet Sun, because y values of the Gaussian fits to the line profiles are high in these areas since the line is very weak there and possibly blended (as suggested by the fact that we obtain mainly strong blueshifts in the umbra contrary to all previous studies based on other spectral lines)."," We have grayed out areas where intensity is below 0.6 times that of the quiet Sun, because $\chi^2$ values of the Gaussian fits to the line profiles are high in these areas since the line is very weak there and possibly blended (as suggested by the fact that we obtain mainly strong blueshifts in the umbra contrary to all previous studies based on other spectral lines)."49" Redshifts (largest value 2.0 ') show a tendency to be located in dark regions, as can be judged by considering the white contours in panel (c)."," Redshifts (largest value 2.0 $^{-1}$ ) show a tendency to be located in dark regions, as can be judged by considering the white contours in panel (c)."50 These contours outline the redshift of panel (d)., These contours outline the redshift of panel (d).51" Redshifts are now found clearly around the head of multiple filaments (e.g.. atv=2"" and y=3"" and around the filaments protruding into the umbra)."," Redshifts are now found clearly around the head of multiple filaments (e.g., at $x=2\varcsec$ and $y=3\varcsec$ and around the filaments protruding into the umbra)."52 Narrow redshifted areas are also found in the middle and outer penumbra beside and between bright filaments., Narrow redshifted areas are also found in the middle and outer penumbra beside and between bright filaments.53" Clearly, in the lower photosphere redshifts are present at many different locations in the penumbra."," Clearly, in the lower photosphere redshifts are present at many different locations in the penumbra."54 We expect that only a part of redshifted features actually present in the penumbra has been detected., We expect that only a part of redshifted features actually present in the penumbra has been detected.55 We have provided the first direct measurements of downflows (reaching 2.0 !) in the body of a penumbra., We have provided the first direct measurements of downflows (reaching 2.0 $^{-1}$ ) in the body of a penumbra.56" The studied sunspot, located at j=0.84, had only a partial penumbra, on the disk-center side of the spot, so that all well-defined filaments partially point to disk center."," The studied sunspot, located at $\mu =0.84$, had only a partial penumbra, on the disk-center side of the spot, so that all well-defined filaments partially point to disk center."57" Hence the Evershed flow contributes a blueshift, so that redshifts in Fig."," Hence the Evershed flow contributes a blueshift, so that redshifts in Fig."58" 5d must be caused by downflows (or inflow, which appears rather unlikely, however)."," 5d must be caused by downflows (or inflow, which appears rather unlikely, however)."59" The line reveals a highly structured velocity pattern, with large variations in velocity around the head"," The line reveals a highly structured velocity pattern, with large variations in velocity around the head"60Radio observations of the 21 cm hydrogen emission. line reveal that the Alilky Way (MW) contains a thin clelisk surrounded by a population of with velocities incompatible with models of galactic rotation vet apparently not part of the Hubble How either (Mulleretal (1963))).,Radio observations of the 21 cm hydrogen emission line reveal that the Milky Way (MW) contains a thin disk surrounded by a population of with velocities incompatible with models of galactic rotation yet apparently not part of the Hubble flow either \cite{mull63}) ).61 The nature of these ‘high-velocity (LIVCs) remains somewhat unclear mostly because it is dillieult to determine their distances and hence infer physical o»operties., The nature of these high-velocity (HVCs) remains somewhat unclear mostly because it is difficult to determine their distances and hence infer physical properties.62 Observations of stars with known distances along the ine of sight to a LIVC can be used to constrain the distance o the cloud. by testing whether or not it is detected in absorption in the stellar spectrum (see e.g. Schwarzctal (1995))).," Observations of stars with known distances along the line of sight to a HVC can be used to constrain the distance to the cloud, by testing whether or not it is detected in absorption in the stellar spectrum (see e.g. \cite{schw95}) )."63 Unfortunately. such constraints. are available for only a relatively small number of LIVCs (Wakker (2001)))., Unfortunately such constraints are available for only a relatively small number of HVCs \cite{wakk01}) ).64 Putmanetal.(2003) use constraints from Lie emission. and find that most LEVC's are within 40 kpe. except for hose associated with the Magellanic stream.," \cite{putm03} use constraints from $\alpha$ emission, and find that most HVCs are within $\sim 40$ kpc, except for those associated with the Magellanic stream."65 Drünsctal. claim that some fraction of LIVCs display a heacl- morphology. may be a result of interaction with a clilfuse ambient. galactic wind (Quilis&Moore (2001))).," \cite{brun01} claim that some fraction of HVCs display a head-tail morphology, may be a result of interaction with a diffuse ambient galactic wind \cite{quil01}) )."66 Searches Or stars associated with Ηλος have so far resulted in non-detections (e.g. Hoppetal. (2007)))., Searches for stars associated with HVCs have so far resulted in non-detections (e.g. \cite{hopp07}) ).67 Metallicities have been measured for a small number of LIVCs and vary over a wide range (e.g. Gibsonetal. (2001))). suggesting that IIVCs are not a homogeneous set.," Metallicities have been measured for a small number of HVCs and vary over a wide range (e.g. \cite{gibs01}) ), suggesting that HVCs are not a homogeneous set."68 Finally one can observe other galaxies to infer the nature of HIVCSs from their projected distances., Finally one can observe other galaxies to infer the nature of HVCs from their projected distances.69 The Andromeda galaxy (M31) has a population of LUVCs close (50 kpe) to its disk (Westmeicr (2007)))., The Andromeda galaxy (M31) has a population of HVCs close $\le 50$ kpc) to its disk \cite{west07}) ).70 Other nearby spiral galaxies also contain neutral gas (see e.g. Barbierictal. (2005)))., Other nearby spiral galaxies also contain extra-planar neutral gas (see e.g. \cite{barb05}) ).71 Pisano(2007) have searched for LIVC analogues in six loose groups of galaxies. similar to the Local Group.," \cite{pisa07} have searched for HVC analogues in six loose groups of galaxies, similar to the Local Group."72 Their failure to detect compact. LIVCs implies that any cclouds are near to the galaxies in this group: they are not roaming [πουν throughout the group. itself., Their failure to detect compact HVCs implies that any clouds are near to the galaxies in this group: they are not roaming freely throughout the group itself.73 These observations do not unambieuously determine the nature of LEVC's. ancl several theoretical models for them have been proposed.," These observations do not unambiguously determine the nature of HVCs, and several theoretical models for them have been proposed."74 Oort (1970) discusses a model where the LWCs are close to the MW. disk. ancl result. from. of a “galactic ((Shapiro&Field (1976))).," Oort (1970) discusses a model where the HVCs are close to the MW disk, and result from of a galactic \cite{shap76}) )."75 Gas from the galactic disk rises buovantly after being heated by supernovae (SNe). becomes thermally unstable and cools raciatively into neutral clouds.," Gas from the galactic disk rises buoyantly after being heated by supernovae (SNe), becomes thermally unstable and cools radiatively into neutral clouds."76 Once the clouds. are dense (μον. fall back ballisticallv onto the disk. and are seen as the λος.," Once the clouds are dense they fall back ballistically onto the disk, and are seen as the HVCs."77 Whether SNe explosions can start a galactic fountain has been investigated both theoretically (Ixahn. (1998))) and through numerical simulations (deAvillez(1998):de&Berry (2001))).," Whether SNe explosions can start a galactic fountain has been investigated both theoretically \cite{kahn98}) ) and through numerical simulations \cite{deav98,deav01}) )."78 Blitz et al. (, Blitz et al. (791999: see also Braun&Burton (1999))) suggest (some) LIVC's are neutral gas associated with the numerous dark matter substructures seen in simulations of haloes of galaxies and groups of galaxies.,1999; see also \cite{brau99}) ) suggest (some) HVCs are neutral gas associated with the numerous dark matter substructures seen in simulations of haloes of galaxies and groups of galaxies.80 Such λος are at large clistances 40 kpe) from the galactic centre. anc," Such HVCs are at large distances $\gtrsim$ 40 kpc) from the galactic centre, and"81in the longitude range 20°</56°. (,in the longitude range $20^\circ <l<56^\circ$. (82Note. however. that there is significant integrated-profile energy in this range and at all longitudes. as shown in the top panel in Fig.,"Note, however, that there is significant integrated-profile energy in this range and at all longitudes, as shown in the top panel in Fig."83 6: this may not be directly related to the drifting subpulses.), 6; this may not be directly related to the drifting subpulses.)84 As can be seen in Fig., As can be seen in Fig.85 3. and more clearly in Fig.," 3, and more clearly in Fig."86 6. the longitude intervals between subpulses (2%) are dillerent in the two prominent regions Land HE. and they vary slightly with longitude.," 6, the longitude intervals between subpulses $P_2$ ) are different in the two prominent regions I and III, and they vary slightly with longitude."87" £5 ranges from LO” to 23.5"" in region Land from 26.8"" to 28 in region HE: in our later discussion we use average values Of 22.2 and 27.5"" for these regions. (", $P_2$ ranges from $19^\circ$ to $23.5^\circ$ in region I and from $26.8^\circ$ to $28^\circ$ in region III; in our later discussion we use average values of $22.2^\circ$ and $27.5^\circ$ for these regions. (88Bs reported that the average £5 in region LIL at 645 MllIz is 297).,B85 reported that the average $P_2$ in region III at 645 MHz is $29^\circ$ ).89 The full width at half maximum (ENIIM) of drifting subpulses in regions I and LL was measured using 20 short sequences of single pulses from both data sets., The full width at half maximum (FWHM) of drifting subpulses in regions I and III was measured using 20 short sequences of single pulses from both data sets.90" Each sequence contained. 20 successive single pulses in a nearly stationary state (Le. drift rate D.«0.5"" in either πι direction).", Each sequence contained 20 successive single pulses in a nearly stationary state (i.e. drift rate $D<0.5^\circ$ in either drift direction).91 For each sequence. the single pulses were co-aclded after adjusting phases of drift to produce one profile.," For each sequence, the single pulses were co-added after adjusting phases of drift to produce one profile."92 The subpulses in region L appear to overlap: we therefore itted Gaussian. curves to. obtain. the ENIM. for each component., The subpulses in region I appear to overlap; we therefore fitted Gaussian curves to obtain the FWHM for each component.93" The ENLEM. of drifting subpulses in region HI ranges rom 10.7"" to 12.77. with an average of 13.17£1.90""."," The FWHM of drifting subpulses in region III ranges from $10.7^\circ$ to $14.7^\circ$, with an average of $13.1^\circ \pm 1.9^\circ$."94 Strong subpulses tend. to. be wider. and to occur in the later ongitudes of the region.," Strong subpulses tend to be wider, and to occur in the later longitudes of the region."95" The ENLIM of subpulses in region I ranges from 16 to 247. with an average of 19d:3.5""."," The FWHM of subpulses in region I ranges from $16^\circ$ to $24^\circ$, with an average of $19^\circ \pm 3.5^\circ$."96 Strong subpulses again end to be wider: they tend to occur around the micelle of he region., Strong subpulses again tend to be wider; they tend to occur around the middle of the region.97 Using the drift phase of each pulse. determined as in Section 4.1. we plotted. the position of cach subpulse in. all. 7600 single pulses of both data sets as shown in Fig.," Using the drift phase of each pulse, determined as in Section 4.1, we plotted the position of each subpulse in all 7600 single pulses of both data sets as shown in Fig."98 7., 7.99 Lhe short sequences of weak-mode pulses. within which the drift phase could not be traced. are indicated by arrows.," The short sequences of weak-mode pulses, within which the drift phase could not be traced, are indicated by arrows."100 The subpulse pattern can drift from. early. to. later longitudes (positive drift) or from late to earlier longitudes (negative drift)., The subpulse pattern can drift from early to later longitudes (positive drift) or from late to earlier longitudes (negative drift).101 The transition from negative to positive drift is usually smooth. while the transition from positive to negative drift is often abrupt.," The transition from negative to positive drift is usually smooth, while the transition from positive to negative drift is often abrupt."102 Phe most rapid drifting. both positive and. negative. occurs immediately before and after," The most rapid drifting, both positive and negative, occurs immediately before and after"103of objects fainter AGN or bright galaxies could. viele uigher precision clustering measurements. perhaps with shotometric samples from surveys such as Pan-STATUS and LSS'T. but perhaps requiring spectroscopic samples like hose envisioned for ambitious barvon acoustic oscillation experiments.,"of objects — fainter AGN or bright galaxies — could yield higher precision clustering measurements, perhaps with photometric samples from surveys such as Pan-STARRS and LSST, but perhaps requiring spectroscopic samples like those envisioned for ambitious baryon acoustic oscillation experiments."104 Phe constraints on host halo populations can also be improved by extending clustering measurements o smaller scales. where quasar pairs from the same halo contribute. and to fainter Iuminosities. such as those probed w the 2dE Quasar Redshift Survey. the SDSS photometric quasar catalog. aud UM SULVOVS (e.&.. llennawi et al.," The constraints on host halo populations can also be improved by extending clustering measurements to smaller scales, where quasar pairs from the same halo contribute, and to fainter luminosities, such as those probed by the 2dF Quasar Redshift Survey, the SDSS photometric quasar catalog, and X-ray surveys (e.g., Hennawi et al."105 2006: Myers et al., 2006; Myers et al.106" 2007:"" Plionis et n", 2007; Plionis et al.107 2008. Llennawi ct al.," 2008, Hennawi et al."108 2009): for example. n et ieOb) use small scale measurements to put mand on duty evele of BIIs in satellite galaxies.," 2009); for example, Shen et (2009b) use small scale measurements to put constraints on the duty cycle of BHs in satellite galaxies."109 Quasar clustering as a cosmological tool has moved from a prospect. (Osmer 1981) to reality. and the growing precision and dynamic. range. of these measurements in Luminosity. redshift. and lengthscale will teach us about— the growth of supermassive black holes and the mechanisms that transform them from. dormant monsters to brilliant beacons. ancl back.," Quasar clustering as a cosmological tool has moved from a prospect (Osmer 1981) to reality, and the growing precision and dynamic range of these measurements — in luminosity, redshift, and lengthscale — will teach us about the growth of supermassive black holes and the mechanisms that transform them from dormant monsters to brilliant beacons, and back."110 FS acknowledges the Alexander von Humboldt Foundation for support., FS acknowledges the Alexander von Humboldt Foundation for support.111 FS and DW also acknowledge support from NASA Grant NNGOSGILIITTCG. and. DW acknowledges support of an AAILAS membership at the Institute. for Advanced: Study.," FS and DW also acknowledge support from NASA Grant NNG05GH77G, and DW acknowledges support of an AMIAS membership at the Institute for Advanced Study."112 We thank Raul Angulo. Luca Craziani. Jeremy Tinker. Jaivul Yoo. and Zheng Zheng for interesting and helpfulcliseussions.," We thank Raul Angulo, Luca Graziani, Jeremy Tinker, Jaiyul Yoo, and Zheng Zheng for interesting and helpful discussions."113 We thank the anonymous referee for comments that led to significant improvements of the paper., We thank the anonymous referee for comments that led to significant improvements of the paper.114not evolve in to a Type la SN as is widely predicted for this system.,not evolve in to a Type Ia SN as is widely predicted for this system.115 Any CN that has been observed to erupt more than once is re-classified as a recurrent nova., Any CN that has been observed to erupt more than once is re-classified as a recurrent nova.116 This approach has given us ten Galactic RNe amongst ~400 known Galactic CNe from an underlying rate of 347 year! (Darnleyetal.2006).., This approach has given us ten Galactic RNe amongst $\sim400$ known Galactic CNe from an underlying rate of $34^{+15}_{-12}$ $^{-1}$ \citep{2006MNRAS.369..257D}.117 Due to observational effects - which typically worsen the further back in time one looks - it i5 highly likely that à considerable proportion of these Galactic CNe are in fact (short time-scale) recurrents for which only a single outburst has been observed., Due to observational effects - which typically worsen the further back in time one looks - it is highly likely that a considerable proportion of these Galactic CNe are in fact (short time-scale) recurrents for which only a single outburst has been observed.118 Either a previous outburst has been missed or the inter-outburst period ts particularly long., Either a previous outburst has been missed or the inter-outburst period is particularly long.119" So it would be seem advantageous to better develop a nova classification system that relies more on the fundamental properties of the systems than on observational selection effects,", So it would be seem advantageous to better develop a nova classification system that relies more on the fundamental properties of the systems than on observational selection effects.120 Over the past few years observations of a number of Galactic “transients” have shown properties consistent with being RNe. although they had never been seen in outburst previously.," Over the past few years observations of a number of Galactic “transients” have shown properties consistent with being RNe, although they had never been seen in outburst previously."121 For example. the nova V2487 Ophiuchi (Nova Oph 1998) was observed to have a rapid optical decline and plateau phase (Hachisuetal.2002;Hernanz&Sala20032). suggestive of a RN.," For example, the nova V2487 Ophiuchi (Nova Oph 1998) was observed to have a rapid optical decline and plateau phase \citep{2002ASPC..261..629H,2002AIPC..637..381H}, suggestive of a RN."122 Consultation of the Harvard College Observatory archival photographie collection revealed a previously unknown outburst on 1900 June 20 2009).. , Consultation of the Harvard College Observatory archival photographic collection revealed a previously unknown outburst on 1900 June 20 \citep{2009AJ....138.1230P}. .123Additionally. the novae V1721 ΑΙ (Hounsell 2011).. V2491 Cyg (Damleyetal.201D... KT Ert (urdana-Sepiéetal.2011) and V2672 Oph (Munarietal.2010) and also the extragalactic nova M31N 2007-12b (Bodeetal.2009) have been shown to harbor secondary stars akin to either the U Sco or RS Oph systems. albeit with only one recorded outburst.," Additionally, the novae V1721 Aql \citep{2011arXiv1104.3068H}, V2491 Cyg \citep{2011arXiv1104.3482D}, KT Eri \citep{RibKTTEri} and V2672 Oph \citep{2010MNRAS.tmp.1484M} and also the extragalactic nova M31N 2007-12b \citep{2009ApJ...705.1056B} have been shown to harbor secondary stars akin to either the U Sco or RS Oph systems, albeit with only one recorded outburst."124 In this paper we will attempt to simplify the nomenclature somewhat by only classifying each system. using. the evolutionary state of the secondary: à main sequence star (MS-Nova). a sub-giant star (SG-Nova). or a red-giant branch star (RG-Nova).," In this paper we will attempt to simplify the nomenclature somewhat by only classifying each system using the evolutionary state of the secondary; a main sequence star (MS-Nova), a sub-giant star (SG-Nova), or a red-giant branch star (RG-Nova)."125 We introduce such classifications in. an attempt to avoid the accretion rate / WD mass / recurrence time degeneracy., We introduce such classifications in an attempt to avoid the accretion rate / WD mass / recurrence time degeneracy.126 All the known U Sco-class RNe would be placed into the SG-Nova group. all the known RS Oph-class RNe into the RG-Nova group. the T Pyx-class RNe and Classical Novae will populate the MS-Nova group.," All the known U Sco-class RNe would be placed into the SG-Nova group, all the known RS Oph-class RNe into the RG-Nova group, the T Pyx-class RNe and Classical Novae will populate the MS-Nova group."127 By virtue of the geometry of these systems. the above classifications would effectively be the same as sub-dividing novae by orbital period. 1f one assumes that the secondaries fill their Roche lobes.," By virtue of the geometry of these systems, the above classifications would effectively be the same as sub-dividing novae by orbital period, if one assumes that the secondaries fill their Roche lobes."128 The MS-Novae would have orbital periods of order hours. the SG-Novae of order a day. and the RG-Novae of order a year.," The MS-Novae would have orbital periods of order hours, the SG-Novae of order a day, and the RG-Novae of order a year."129 Indeed Warner(1995) indicates that systems with orbital periods longer than eight. hours should contain evolved secondaries., Indeed \citet{1995CAS....28.....W} indicates that systems with orbital periods longer than eight hours should contain evolved secondaries.130 In this paper we present a simple method which will allow differentiation of RG-Novae and SG-Novae from the MS-Novae population based solely on their quiescent (inter-outburst) broad-band optical and near-IR (NIR) properties., In this paper we present a simple method which will allow differentiation of RG-Novae and SG-Novae from the MS-Novae population based solely on their quiescent (inter-outburst) broad-band optical and near-IR (NIR) properties.131 The format of the remainder of this paper/letter is as follows: In Section 2? we briefly describe our method and present the data used to illustrate this technique: in Section 22. we present our results. in Section ?? we discuss our findings and present a number of predictions and finally. in Section ??. we summarize our conclusions.," The format of the remainder of this paper/letter is as follows: In Section \ref{data} we briefly describe our method and present the data used to illustrate this technique; in Section \ref{results} we present our results, in Section \ref{discussion} we discuss our findings and present a number of predictions and finally, in Section \ref{conc} we summarize our conclusions."132 The broad-band optical emission from all quiescent nova systems I8 a super-position of the emission from the three main components of the system: the WD. the accretion disk. and the secondary.," The broad-band optical emission from all quiescent nova systems is a super-position of the emission from the three main components of the system; the WD, the accretion disk, and the secondary."133 In all cases we would expect the WD's contribution m the optical to be negligible., In all cases we would expect the WD's contribution in the optical to be negligible.134 The contribution of the accretion disk to the emission is a combination of a number of factors including: accretion rate. disk size. system inclination and wavelength.," The contribution of the accretion disk to the emission is a combination of a number of factors including; accretion rate, disk size, system inclination and wavelength."135 Whereas the contribution from the secondary is much more straightforward and simply depends upon the type (mass. age and metalicity) of the star and wavelength.," Whereas the contribution from the secondary is much more straightforward and simply depends upon the type (mass, age and metalicity) of the star and wavelength."136 Both the SG-Novae and RG-Novae have much smaller optical outburst amplitudes than MS-Novae. although the energetics of the outbursts of all three systems are comparable.," Both the SG-Novae and RG-Novae have much smaller optical outburst amplitudes than MS-Novae, although the energetics of the outbursts of all three systems are comparable."137 In the RG-Nova systems this can be readily explained by the presence of a - highly luminous - red giant secondary. possibly with some additional contribution from a disk. especially in bluer optical filters.," In the RG-Nova systems this can be readily explained by the presence of a - highly luminous - red giant secondary, possibly with some additional contribution from a disk, especially in bluer optical filters."138 For the SG-Nova systems the cause is not as clear-cut but must be some combination of a - luminous - sub-giant secondary and an accretion disk. possibly experience a high mass accretion rate.," For the SG-Nova systems the cause is not as clear-cut but must be some combination of a - luminous - sub-giant secondary and an accretion disk, possibly experience a high mass accretion rate."139 The broad-band NIR emission from quiescent RG-Nova and SG-Nova systems should be dominated by the emission from the cool secondary stars. with a small contributioαυ] from the aceretion disk.," The broad-band NIR emission from quiescent RG-Nova and SG-Nova systems should be dominated by the emission from the cool secondary stars, with a small contribution from the accretion disk."140 At these wavelengths. the accretion disk emission may still be relatively significant for MS-Nova systems due to the low luminosity of the dwarf secondary star.," At these wavelengths, the accretion disk emission may still be relatively significant for MS-Nova systems due to the low luminosity of the dwarf secondary star."141 As such. we propose that the position of a quiescer= nova system on a color-magnitude diagram can be broadly predicted.," As such, we propose that the position of a quiescent nova system on a color-magnitude diagram can be broadly predicted."142 Each system would be expected to lie at such a position that it appears more luminous and bluer (hotter) than the secondary in the system. with the magnitude of the luminosity and temperature increase being a function of both the accretion rate. system inclination and - of course - wavelength.," Each system would be expected to lie at such a position that it appears more luminous and bluer (hotter) than the secondary in the system, with the magnitude of the luminosity and temperature increase being a function of both the accretion rate, system inclination and - of course - wavelength."143 Conversely. by plotting quiescent nova systems on a color-magnitude diagram one can identify the type of secondary and hence predict whether a given system may be an RG-Nova. SG-Nova or MS-Nova system.," Conversely, by plotting quiescent nova systems on a color-magnitude diagram one can identify the type of secondary and hence predict whether a given system may be an RG-Nova, SG-Nova or MS-Nova system."144 In Table 2. we present quiescent broad-band optical and NIR photometric data for a sample of Galactic novae (and a number of extragalactic examples)., In Table \ref{tb:one} we present quiescent broad-band optical and NIR photometric data for a sample of Galactic novae (and a number of extragalactic examples).145 These data have been collated from a wide range of sources. with all the NIR data taken from the Two Micron All Sky Survey (2MASS:Skrut-skieetal.2006) either directly from the catalog or indirectly through additional sources.," These data have been collated from a wide range of sources, with all the NIR data taken from the Two Micron All Sky Survey \citep[2MASS;][]{2006AJ....131.1163S} either directly from the catalog or indirectly through additional sources."146 In Table | we provide a summary of the best determination of the line of sight extinction and distance to each system as well as the system inclination (if known)., In Table \ref{tb:two} we provide a summary of the best determination of the line of sight extinction and distance to each system as well as the system inclination (if known).147 This catalog of Galactic novae was compiled by virtue of the existence of quiescent photometry in at least two optical bands. and reasonable determinations of both distance and extinction.," This catalog of Galactic novae was compiled by virtue of the existence of quiescent photometry in at least two optical bands, and reasonable determinations of both distance and extinction."148 In Figure we present a range of color-magnitude plots., In Figure \ref{fig:one} we present a range of color-magnitude plots.149 The positions1. of all the novae listed in Table 2 are plotted with appropriateerror bars (or error estimations)., The positions of all the novae listed in Table \ref{tb:one} are plotted with appropriateerror bars (or error estimations).150 For comparison. we have also plotted the local Galactic stellar population. generated from the catalog (Perryman 1997)..," For comparison, we have also plotted the local Galactic stellar population, generated from the catalog \citep{1997ESASP1200.....P}."151 Only those stars that have both photometric and parallax errors <10% have been selected from theHipparcos data., Only those stars that have both photometric and parallax errors $<10\%$ have been selected from the data.152 The BVI photometry for these stars has been taken directly from that catalog. while the R and NIRphotometry has been generated by cross-correlating the catalog with the data set (Zacharias and the 2MASS cataogue respectively.," The $BVI$ photometry for these stars has been taken directly from that catalog, while the $R$ and NIRphotometry has been generated by cross-correlating the catalog with the data set \citep{2004AAS...205.4815Z} and the 2MASS cataogue respectively."153 Included, Included154"that from compression, and about equal to that from formation.","that from compression, and about equal to that from formation."155" We calculate the heating rate from compressional heating using the below formula that can be derived from energy conservation where € is the thermal energy per unit volume, y is the adiabatic index of the gas and v is the velocity."," We calculate the heating rate from compressional heating using the below formula that can be derived from energy conservation where $\epsilon$ is the thermal energy per unit volume, $\gamma$ is the adiabatic index of the gas and v is the velocity."156" The heating from compression and formation was already being balanced by cooling from line emission, and re-expansion of the gas, as shown in 8.."," The heating from compression and formation was already being balanced by cooling from line emission, and re-expansion of the gas, as shown in \ref{cooling}."157" Therefore, the addition of accretion luminosity feedback represents only a small change in the thermodynamic equilibrium of the halo."," Therefore, the addition of accretion luminosity feedback represents only a small change in the thermodynamic equilibrium of the halo."158" However, an important qualification is that luminosity feedback is an effect that varies with position, i.e. it is most effective close to the sinks where the gas is densest and fragmentation occurs."," However, an important qualification is that luminosity feedback is an effect that varies with position, i.e. it is most effective close to the sinks where the gas is densest and fragmentation occurs."159" Additionally, the extra heating increases the collisional dissociation rate ofΗ2., making it harder for the dense gas to cool."," Additionally, the extra heating increases the collisional dissociation rate of, making it harder for the dense gas to cool."160" Consequently, the accretion luminosity heating is more dynamically significant than a first glance at 7 would suggest, which is why it was able to delay fragmentation in the minihalos."," Consequently, the accretion luminosity heating is more dynamically significant than a first glance at \ref{heating} would suggest, which is why it was able to delay fragmentation in the minihalos."161 The previous section demonstrates that accretion luminosity can affect the fragmentation seen in the minihalos., The previous section demonstrates that accretion luminosity can affect the fragmentation seen in the minihalos.162 However it has also shown that these effects can be completely masked by the dynamics of the gas., However it has also shown that these effects can be completely masked by the dynamics of the gas.163" A similar conclusion is derived from massive star formation calculations in the present day, where dynamical effects dominate over radiative feedback (??,a,b).."," A similar conclusion is derived from massive star formation calculations in the present day, where dynamical effects dominate over radiative feedback \citep[][a,b]{Krumholz09,Peters10b}."164" To explore this more fully, in this section we contrast dynamical effects with those of accretion luminosity upon the individual protostars, and consider how this will affect our feedback model."," To explore this more fully, in this section we contrast dynamical effects with those of accretion luminosity upon the individual protostars, and consider how this will affect our feedback model."165 9 shows the evolution of the sinks formed in halos 1-4., \ref{sinks} shows the evolution of the sinks formed in halos 1-4.166 The first sink forms at the centre of the disk and quickly grows in mass with a smoothly decreasing accretion rate., The first sink forms at the centre of the disk and quickly grows in mass with a smoothly decreasing accretion rate.167" ? showed that the expected accretion rates should be as high as 107! after the first half solar mass has collapsed, and then smoothly decrease to a value of of 107? when 100 has collapsed."," \citet{Yoshida06} showed that the expected accretion rates should be as high as $10^{-1}$ after the first half solar mass has collapsed, and then smoothly decrease to a value of of $10^{-3}$ when $100$ has collapsed."168 Our accretion rate for the initial sink agrees with ? until fragmentation sets in., Our accretion rate for the initial sink agrees with \citet{Yoshida06} until fragmentation sets in.169" At this point the accretion rate may briefly rise as the portion of the disk between the original sink and the new sinks is strongly torqued, resulting in a large outward transfer of angular momentum and inflow of gas onto the central sink."," At this point the accretion rate may briefly rise as the portion of the disk between the original sink and the new sinks is strongly torqued, resulting in a large outward transfer of angular momentum and inflow of gas onto the central sink."170" After this short transient, the accretion rate decreases as the mass available for accretion is now shared between multiple sinks."," After this short transient, the accretion rate decreases as the mass available for accretion is now shared between multiple sinks."171" Once multiple sinks are formed, the accretion rates of the sink particles become highly variable."," Once multiple sinks are formed, the accretion rates of the sink particles become highly variable."172" Due to the high densities characteristic of primordial star formation, the Jeans length is extremely short, and therefore fragments are formed close to each other, leading to interactions on timescales comparable to the local free-fall time."," Due to the high densities characteristic of primordial star formation, the Jeans length is extremely short, and therefore fragments are formed close to each other, leading to interactions on timescales comparable to the local free-fall time."173 9 also shows the paths of the stars in the central 4000 AU during the period studied here., \ref{sinks} also shows the paths of the stars in the central $4000$ AU during the period studied here.174" The sinks that remain in the halo centre orbit each other, leading to a periodically varying accretion rate as they move around in the gas bound to them."," The sinks that remain in the halo centre orbit each other, leading to a periodically varying accretion rate as they move around in the gas bound to them."175" Moreover, ejections are common and occur in every halo, as shown by ?.."," Moreover, ejections are common and occur in every halo, as shown by \citet{Greif11}."176" Surprisingly, even the originally central star can be ejected if it has a close three-body interaction as shown in Halo 1 (this only occurred in the case with feedback, which is why it was hard to compare Halo 1 with and without feedback)."," Surprisingly, even the originally central star can be ejected if it has a close three-body interaction as shown in Halo 1 (this only occurred in the case with feedback, which is why it was hard to compare Halo 1 with and without feedback)."177" Indeed, while we find that feedback from accretion luminosity has no significant effect on the accretion rates of the protostars, dynamical interactions are extremely effective at halting an ejected protostar's accretion entirely (??).."," Indeed, while we find that feedback from accretion luminosity has no significant effect on the accretion rates of the protostars, dynamical interactions are extremely effective at halting an ejected protostar's accretion entirely \citep{Reipurth01,Bate02}."178" The accretion rates for the sink particles are variable, yet the stellar radius model used to calculate the accretion luminosity was developed from simulations with a constant accretion rate."," The accretion rates for the sink particles are variable, yet the stellar radius model used to calculate the accretion luminosity was developed from simulations with a constant accretion rate."179 Figure 10 shows the stellar radius that results from the measured sink accretion rates using the stellar radius model shown in 1.., Figure \ref{sinkradius} shows the stellar radius that results from the measured sink accretion rates using the stellar radius model shown in \ref{radmodel}.180 The expected trend of an increasing radius which reaches a sharp peak and then rapidly decreases is still found., The expected trend of an increasing radius which reaches a sharp peak and then rapidly decreases is still found.181" However, variation"," However, variation"182 Recent attempts to model the spectiuui of Sagittarius A* have proven promising. but have relied upon siuplifications or assuniptiouns which have somewhat diminished the impact of these results.," Recent attempts to model the spectrum of Sagittarius A* have proven promising, but have relied upon simplifications or assumptions which have somewhat diminished the impact of these results."183 Our intent is o deliver the most consistent modeling possible with current technology. and avoid some of the simplification hat previous simulations have required.," Our intent is to deliver the most consistent modeling possible with current technology, and avoid some of the simplifications that previous simulations have required."184 This begius with using a GRATID code specifically designed to evolve an accretion disk around a black role., This begins with using a GRMHD code specifically designed to evolve an accretion disk around a black hole.185 This is a considerable advantage over models which specify analytic formulae to describe deusity. cluperature. aud magnetic field profiles.," This is a considerable advantage over models which specify analytic formulae to describe density, temperature, and magnetic field profiles."186 Even simulations done with MIID codes suffer some of the sae downfalls as slunpler set-üps specifically that effects of eeucral relativity near à supermassive black hole become quite nuportaut., Even simulations done with MHD codes suffer some of the same downfalls as simpler set-ups — specifically that effects of general relativity near a supermassive black hole become quite important.187 Data output bv the GRMIID code are used as nupt for our MC code., Data output by the GRMHD code are used as input for our MC code.188 Uulike other models which rely ouly on emission aud scattering calculated aualvticallv. this sinulation creates a realistic situation through use of the Monte-Carlo random generation method for photous. and allows for scattering aud absorption before photous escape the volume.," Unlike other models which rely only on emission and scattering calculated analytically, this simulation creates a realistic situation through use of the Monte-Carlo random generation method for photons, and allows for scattering and absorption before photons escape the volume."189 It should be noted that both of the codes used here ave 2D and axialle svaauuetries, It should be noted that both of the codes used here are 2D and axially symmetric.190 The third dimension would be useful for full consistency. but the authors feel that its lack of inclusion docs not preclude the results from being significant.," The third dimension would be useful for full consistency, but the authors feel that its lack of inclusion does not preclude the results from being significant."191 Ow GRMIID inodeliue is done using the TARNI (igh Accuracy Relativistic Maeuctolydrocdwuamics) code developed. by CGamunieetal.(2003)., Our GRMHD modeling is done using the HARM (High Accuracy Relativistic Magnetohydrodynamics) code developed by \cite{gam03}.192. The code siuulates GRATID evolution through time by the use of particle nuuber couservation. the four cucrey-momentum equations. the MIID. stress-cuerey tensor.," The code simulates GRMHD evolution through time by the use of particle number conservation, the four energy-momentum equations, the MHD stress-energy tensor,"193the white dwarf temperature.,the white dwarf temperature.194" The spot temperature is subject to much larger systematic uncertainties, since e.g. our assumption of a uniform spot temperature is a very crude approximation."," The spot temperature is subject to much larger systematic uncertainties, since e.g. our assumption of a uniform spot temperature is a very crude approximation."195 The model spectra shown in Fig., The model spectra shown in Fig.196" 3 at orbital minimum and maximum were computed for a binary inclication i=60°, a spot extent of 24? (half opening angle), and a spot colatitude of just 12.5?."," \ref{f:seduvopt}197 at orbital minimum and maximum were computed for a binary inclication $i=60\degr$, a spot extent of $24\degr$ (half opening angle), and a spot colatitude of just $12.5\degr$."198" A rather high inclination and a high ""northern? spot latitude, the spot undergoes a partial self-eclipse only, are required from the fact that the FUV-band is almost completely dominated by the spot."," A rather high inclination and a high 'northern' spot latitude, the spot undergoes a partial self-eclipse only, are required from the fact that the FUV-band is almost completely dominated by the spot."199 Even at orbital minimum the white dwarf contributes only ~10% to the total flux in that band (see the lowest model curve in Fig. 3))., Even at orbital minimum the white dwarf contributes only $\sim$ to the total flux in that band (see the lowest model curve in Fig. \ref{f:seduvopt}) ).200 Our model is simple and far from being unique but it fits the data well and contradicts the conclusion by Szkody et al. (, Our model is simple and far from being unique but it fits the data well and contradicts the conclusion by Szkody et al. (2012006) that no spot model can explain both the SED and the variability.,2006) that no spot model can explain both the SED and the variability.202" Combining parallaxes, proper motions and absolute magnitude constraints, Thorstensen (2003) derived distance estimates for 14 CVs with a Bayesian method, among themEri."," Combining parallaxes, proper motions and absolute magnitude constraints, Thorstensen (2003) derived distance estimates for 14 CVs with a Bayesian method, among them."203" Dependent on the proper-motion, magnitude and velocity priors, he derived a short, d—113*7 ppc, and a long, d=163796 ppc, distance toEFEri."," Dependent on the proper-motion, magnitude and velocity priors, he derived a short, $d = 113^{+19}_{-16}$ pc, and a long, $d =204163^{+66}_{-50}$ pc, distance to."205. At ppc the observed flux of our 9750 K white dwarf model results in a radius of 6.5x108 ccm and implies a relatively massive white dwarf ofmsun., At pc the observed flux of our 9750 K white dwarf model results in a radius of $6.5 \times 10^{8}$ cm and implies a relatively massive white dwarf of.206. At ppc the implied mass is ~0.55msun., At pc the implied mass is $\sim$.207. These numbers differ slightly from those derived in Beuermann et al. (, These numbers differ slightly from those derived in Beuermann et al. (208"2000), since the spot contribution was taken into account in our analysis.","2000), since the spot contribution was taken into account in our analysis."209 We have analysed archival XMM-Newton observations of oobtained in 2002 and 2003., We have analysed archival XMM-Newton observations of obtained in 2002 and 2003.210" At all three occasions the polar was detected as an X-ray source, although at a very low flux level."," At all three occasions the polar was detected as an X-ray source, although at a very low flux level."211 The spectra were compatible with emission from a low-temperature corona-like plasma., The spectra were compatible with emission from a low-temperature corona-like plasma.212" The longest X-ray observation had almost full phase coverage, the plasma temperature was as low as kkeV or less."," The longest X-ray observation had almost full phase coverage, the plasma temperature was as low as keV or less."213 The mean orbital integrated flux in this component is about Fy=6x107'ergs!..," The mean orbital integrated flux in this component is about $F_X =2146 \times 10^{-14}$."215" Assuming isotropic radiation and a distance of ppc (Thorstensen 2003), a luminosity of Ly~2x10°9 is derived."," Assuming isotropic radiation and a distance of pc (Thorstensen 2003), a luminosity of $L_X \sim 2 \times 10^{29}$ is derived."216 The question arises if this faint X-ray flux originates from the corona of the secondary or from the accretion region on the white dwarf., The question arises if this faint X-ray flux originates from the corona of the secondary or from the accretion region on the white dwarf.217 Neither X-ray variability nor the X-ray spectrum give a clear answer., Neither X-ray variability nor the X-ray spectrum give a clear answer.218" Both, a coronal plasma and the cooling plasma from low level accretion have such low temperatures as measured here."," Both, a coronal plasma and the cooling plasma from low level accretion have such low temperatures as measured here."219 Evidence for X-ray emission from an accretion plasma can be given indirectly., Evidence for X-ray emission from an accretion plasma can be given indirectly.220" Firstly, although not very much is known about X-ray emission from degenerate stars at the bottom of the main sequence, their X-ray luminosities seem to fall short by one dex with respect to the X-ray luminosiy of (Stelzer et al."," Firstly, although not very much is known about X-ray emission from degenerate stars at the bottom of the main sequence, their X-ray luminosities seem to fall short by one dex with respect to the X-ray luminosiy of (Stelzer et al."221 2006)., 2006).222" Secondly, sshows clear signs of residual accretion via the detection of infrared cyclotron harmonics (Harrison et al."," Secondly, shows clear signs of residual accretion via the detection of infrared cyclotron harmonics (Harrison et al."223 2004)., 2004).224 It therefore appears reasonable to assign the observed X-ray emission to some remaining weak accretion., It therefore appears reasonable to assign the observed X-ray emission to some remaining weak accretion.225 This will be our working hypothesis in the following., This will be our working hypothesis in the following.226" It remains unclear if residual accretion happens via an accretion stream or via a stellar wind, the latter being inferred in order to explain the faint X- emission from the small group of pre-CVs (termed also LARPs, Schwope et al."," It remains unclear if residual accretion happens via an accretion stream or via a stellar wind, the latter being inferred in order to explain the faint X-ray emission from the small group of pre-CVs (termed also LARPs, Schwope et al."227" 2002, Schmidt et al."," 2002, Schmidt et al."228" 2005, Vogel et al."," 2005, Vogel et al."229 2006)., 2006).230 Since the emission region in iis not self-eclipsing (Beuermann et al., Since the emission region in is not self-eclipsing (Beuermann et al.231" 1987, 1991), we are lacking a distinct photometric feature to discern between the two possibilities."," 1987, 1991), we are lacking a distinct photometric feature to discern between the two possibilities."232" The pronounced soft X-ray absorption dip as a sign of stream accretion and seen in high accretion states was not secularly detected here, but due to the small number of photons its absence does not give a clear-cut answer to the question of the accretion mode."," The pronounced soft X-ray absorption dip as a sign of stream accretion and seen in high accretion states was not secularly detected here, but due to the small number of photons its absence does not give a clear-cut answer to the question of the accretion mode."233 We discuss the energy balance of the accretion process in the low state on the assumption that the observed X-ray emission is due to accretion onto the white dwarf primary., We discuss the energy balance of the accretion process in the low state on the assumption that the observed X-ray emission is due to accretion onto the white dwarf primary.234 We make the further assumption that the excess emission in the infrared over the extrapolated white dwarf spectrum is solely due to cyclotron emission from the accretion plasma., We make the further assumption that the excess emission in the infrared over the extrapolated white dwarf spectrum is solely due to cyclotron emission from the accretion plasma.235 The relevant radiation components are shown in Fig. 4.., The relevant radiation components are shown in Fig. \ref{f:sedirx}.236" It shows the coronal plasma in the X-ray regime and the cyclotron component in the infrared, the latter corrected for the contribution from the underlying white dwarf."," It shows the coronal plasma in the X-ray regime and the cyclotron component in the infrared, the latter corrected for the contribution from the underlying white dwarf."237" Included in the figure is the spot model at orbital maximum, represented by a white dwarf model spectrum with Teg=18500 KK."," Included in the figure is the spot model at orbital maximum, represented by a white dwarf model spectrum with $\teff = 18500$ K."238rms-ILux relation (which applies to the entire light curve) in case (i). but the average variance is the same in both cases.,"rms-flux relation (which applies to the entire light curve) in case (i), but the average variance is the same in both cases."239 In Fig., In Fig.240 3. we plot the observed. pulse rms. versus aperiocdic variance., \ref{obsrmsvar} we plot the observed pulse rms versus aperiodic variance.241 The pulse rms and the aperiodic variance are clearly correlatec eracicnt (2.2+0.4)10 for yO2= 3.2. for 2 degrees of freedom).," The pulse rms and the aperiodic variance are clearly correlated (gradient $(2.2\pm0.4)\times10^{-4}$ for $\chi^{2}=3.2$ , for 2 degrees of freedom)."242 Furthermore. the observed eracient of the correlation Is consistent with that predictecl ov the case (i) simulated: cata.," Furthermore, the observed gradient of the correlation is consistent with that predicted by the case (i) simulated data."243 Pherefore. we conclude hat the aperiodic variability containing the linear rnis-lux relation is coupled to the pulse. and. furthermore. no additional component of aperiodic variability is required by he data.," Therefore, we conclude that the aperiodic variability containing the linear rms-flux relation is coupled to the pulse, and furthermore, no additional component of aperiodic variability is required by the data."244 Thus we infer that the X-ray variability carrying he rms-lux relation originates> at the magnetico caps of he neutron star and therefore the P?> model is. correct in SAX J1808.4-3658 and the CE model is ruled out., Thus we infer that the X-ray variability carrying the rms-flux relation originates at the magnetic caps of the neutron star and therefore the PP model is correct in SAX J1808.4-3658 and the CF model is ruled out.245 For completeness. we note here the possibility that. case (i) corresponds to a more complicated PP model. where most of the aperiodic variability is produced. before. the »ulsed: variability. by perturbations propagating through a separate unpulsed emitting region (e.g. a static corona above he disk). before reaching the site of the pulsed. emission.," For completeness, we note here the possibility that case (i) corresponds to a more complicated PP model, where most of the aperiodic variability is produced before the pulsed variability by perturbations propagating through a separate unpulsed emitting region (e.g. a static corona above the disk), before reaching the site of the pulsed emission."246 Llowever. given the available spectral-variability. evidence ‘or other neutron stars (Cillanov.Revnivisey& 2003).. it is simplest to assume that all the variable emission originates at or close to the neutron star surface.," However, given the available spectral-variability evidence for other neutron stars \citep{gil03}, it is simplest to assume that all the variable emission originates at or close to the neutron star surface."247 We have demonstrated that the aperiodic X-ray. variations which contain the linear rnis-IHux relation in SAX J180S.4-3658 are coupled to the 401 Hz pulsation. and hence originate on or close to the magnetic caps of the neutron star.," We have demonstrated that the aperiodic X-ray variations which contain the linear rms-flux relation in SAX J1808.4-3658 are coupled to the 401 Hz pulsation, and hence originate on or close to the magnetic caps of the neutron star."248 The only existing mocel which can explain this result is the PP model: specifically. the rmis-IHux relation is produced by the coupling of perturbations in the aceretion [low as they propagate inwards.," The only existing model which can explain this result is the PP model: specifically, the rms-flux relation is produced by the coupling of perturbations in the accretion flow as they propagate inwards."249 When the accretion flow is channeled on to the magnetic caps of the neutron star. the energy of accretion is released and the pattern of Ductuations in the accretion How (including the rms-Iux relation) is imprinted on the resulting X-ray emission.," When the accretion flow is channeled on to the magnetic caps of the neutron star, the energy of accretion is released and the pattern of fluctuations in the accretion flow (including the rms-flux relation) is imprinted on the resulting X-ray emission."250 Although we have only shown that the PP model is the most likely explanation of the rms-Iux relation in SAN JISOS.4-3658. it seems highly likely also that the PP model explains the aperiodic variability and. rmis-Hux. relation in other accreting neutron star ancl black hole svstems. (including AGN).," Although we have only shown that the PP model is the most likely explanation of the rms-flux relation in SAX J1808.4-3658, it seems highly likely also that the PP model explains the aperiodic variability and rms-flux relation in other accreting neutron star and black hole systems (including AGN)."251 For example. the power spectra of broadband: noise in both black hole and neutron star NRBs (including SAN JI1808.4-3658) can be described by a simple model involving the superposition of broad. Lorentzian features (Belloni.Psaltis&vanderIxlis2002:Pottschmicltetal.2003:vanStraaten.derIxlis&Wijnands 2003)). and both neutron star and. black hole systems show similar correlations between the dillerent characteristic frequencies of these features (c.g. Wijnands&vancerExlis1999:Bel-loni.Psaltis&vanderIxlis2002:Pottschimidtetal. 2003)).," For example, the power spectra of broadband noise in both black hole and neutron star XRBs (including SAX J1808.4-3658) can be described by a simple model involving the superposition of broad Lorentzian features \citealt{bel02,pot03,str03}) ), and both neutron star and black hole systems show similar correlations between the different characteristic frequencies of these features (e.g. \citealt{wij99,bel02,pot03}) )."252 The strong similarities between black hole ΧΙ variability and that of AGN. with characteristic time-scales apparentIv scaling with the black hole mass (c.g. Uttlev.MLlardy&Pa-padakis2002:Markowitzetal.2003:Al’Lares 2003)) also strongly support the idea that AGN and XItDs share the same aperiodic variability mechanism.," The strong similarities between black hole XRB variability and that of AGN, with characteristic time-scales apparently scaling with the black hole mass (e.g. \citealt{utt02,mar03,mch03}) ) also strongly support the idea that AGN and XRBs share the same aperiodic variability mechanism."253 Since the light curves of all these tvpes of system. show a linear rms-Iux relation (c.g. Uttlev&M'LEbardy2001:Vaughan.FabianNandra2003:Gleissnerοἱal. 2003)). the application of Ockham's tazor would suggest that the rms-Iux relation is produces w the same mechanism in all cases and hence the aperiocdic X-ray variabilitv in neutron star and black bole NliDs aux AGN is produced by a PP mechanism. and not by corona lares.," Since the light curves of all these types of system show a linear rms-flux relation (e.g. \citealt{utt01,vau03,gle03}) ), the application of Ockham's Razor would suggest that the rms-flux relation is produced by the same mechanism in all cases and hence the aperiodic X-ray variability in neutron star and black hole XRBs and AGN is produced by a PP mechanism, and not by coronal flares."254 Lt should. be restated here that the PP interpretation of he variability does not carry any implications for the existence of an X-ray emitting corona., It should be restated here that the PP interpretation of the variability does not carry any implications for the existence of an X-ray emitting corona.255 This is because the aperiodic variability (and. the rms-dux. relation. imprintec in it) is produced in the accretion Low and is independen of where the X-rays are emitted. provided that. variations in the accretion [low can modulate the X-ray. emission.," This is because the aperiodic variability (and the rms-flux relation imprinted in it) is produced in the accretion flow and is independent of where the X-rays are emitted, provided that variations in the accretion flow can modulate the X-ray emission."256 Thus a corona which is heated by many small reconnection events (too small to have a very strong ellect on. the variability) remains a viable source of the X-ray emission (in black hole svstems. at least).," Thus a corona which is heated by many small reconnection events (too small to have a very strong effect on the variability) remains a viable source of the X-ray emission (in black hole systems, at least)."257 In fact. by invoking an extended. X-ray. emitting region (such as à corona) which »ossesses à temperature eradient. so that higher energy. ravs are prelerentially emitted closer to the black hole. PP- models naturally produce the time-dependent: delays oween energy. bands and the energy-dependent. shape of le power spectrum. which are observed in black hole ray binaries (c.g. Alisra2000:Ixotov.2001:Zvcki 2003)) and AGN (e.g. dra2003:MLlardy.ctal. 2003)).," In fact, by invoking an extended X-ray emitting region (such as a corona) which possesses a temperature gradient, so that higher energy X-rays are preferentially emitted closer to the black hole, PP-type models naturally produce the time-dependent delays between energy bands and the energy-dependent shape of the power spectrum, which are observed in black hole X-ray binaries (e.g. \citealt{mis00,kot01,zyc03}) ) and AGN (e.g. \citealt{vau03,mch03}) )."258 We have recently shown τον.AlUardy&Vaughan 2003.. ancl Uttlev et al.," We have recently shown \citealt{utt03}, and Uttley et al.,"259 in prep.), in prep.)260 that the existence of an rms-Hux relation on all time-scales leads naturally to the appearance of non-linear behaviour which is observed. over a range of time-scales in both NRBs (Alaccarone 2003)) and AGN," that the existence of an rms-flux relation on all time-scales leads naturally to the appearance of non-linear behaviour which is observed over a range of time-scales in both XRBs \citealt{mac02,gie03}) ) and AGN"261"The nullsets in the preceding section may naturally be replaced by an abstract notion of ""negligible"" sets, i.e, members of a (briefly: p.131.","The nullsets in the preceding section may naturally be replaced by an abstract notion of “negligible” sets, i.e. members of a (briefly: p.l.i."262 ideal) in E. defined as follows., ideal) in $\R$ defined as follows.263 A nonempty family 7C2-\[IE] is termed to be a p.Li., A nonempty family $\id\subset 2^{\R} \setminus \{\R\}$ is termed to be a p.l.i.264 ideal (resp., ideal (resp.265 p.l.i., p.l.i.266 —ideal) provided that it is closed under finite (resp., $\sigma-$ ideal) provided that it is closed under finite (resp.267" countable) set theoretical unions, i.e. hereditary with respect to descending inclusions, i.e. and such that jointly with a given set it contains its image under any alfine transformation of the real line onto itself, i.e Clearly the family of all nullsets (sets of Lebesgue measure zero) in IE. forms a p.l.i."," countable) set theoretical unions, i.e. hereditary with respect to descending inclusions, i.e. and such that jointly with a given set it contains its image under any affine transformation of the real line onto itself, i.e Clearly the family of all nullsets (sets of Lebesgue measure zero) in $\R$ forms a p.l.i."268 o-ideal., $\sigma$ -ideal.269 However there are numerous other p.l.i., However there are numerous other p.l.i.270 ideals: let us mention only a few of them:, ideals; let us mention only a few of them:271eemission-line flux for 26/34 (~ 75%)) galaxies in the sample.,emission-line flux for 26/34 $\sim 75$ ) galaxies in the sample.272 The remaining 8 galaxies have either a too low SNR or the eemission-line is polluted by sky-line residuals., The remaining 8 galaxies have either a too low SNR or the emission-line is polluted by sky-line residuals.273 A recent study (?) has shown that unveiling the presence of AGN in high-redshift galaxies is a difficult exercise., A recent study \citep{wright10} has shown that unveiling the presence of AGN in high-redshift galaxies is a difficult exercise.274" In the case of metallicity studies, in which abundances are deduced from the ratio of different emission lines of the ionised gas, it is critical to check that the intensity and width of these lines are due to star formation and not related to any non-thermal nuclear activity."," In the case of metallicity studies, in which abundances are deduced from the ratio of different emission lines of the ionised gas, it is critical to check that the intensity and width of these lines are due to star formation and not related to any non-thermal nuclear activity."275" The common way to disentangle AGN contribution from star-forming galaxies consists in comparing the relative intensity of the main nebular emission lines (mainly textsciii]]|5007, H8, Ha, and [N1]6584)) in a diagnostic diagram, so-called BPT diagram (??)."," The common way to disentangle AGN contribution from star-forming galaxies consists in comparing the relative intensity of the main nebular emission lines (mainly ]5007, $\beta$, $\alpha$, and ) in a diagnostic diagram, so-called BPT diagram \citep{bpt81, kewley01}."276" Various physical conditions in the ISM -— SFR, ionisation parameter, metallicity and/or chemical composition -— have been invoked to explain the fact that some high-z star-forming galaxies lie in the transition region of the local BPT diagram (as defined by the SDSS galaxies) between star-forming galaxies and AGN hosts."," Various physical conditions in the ISM – SFR, ionisation parameter, metallicity and/or chemical composition – have been invoked to explain the fact that some $z$ star-forming galaxies lie in the transition region of the local BPT diagram (as defined by the SDSS galaxies) between star-forming galaxies and AGN hosts."277" ? have been able -— thanks to high resolution adaptive optics observations -— to subtract the active nuclear emission in a z~1.6 galaxy (HDF-BMZ1299), and have shown that the residual extended star-forming emission was characteristic of a local SDSS star-forming galaxy, whereas the integrated emission would have placed the object in the transition region."," \citet{wright10} have been able – thanks to high resolution adaptive optics observations – to subtract the active nuclear emission in a $z\sim 1.6$ galaxy (HDF-BMZ1299), and have shown that the residual extended star-forming emission was characteristic of a local SDSS star-forming galaxy, whereas the integrated emission would have placed the object in the transition region."278" The nature of our observations did not give us simultaneous access to the set of emission lines textsciii]]5007, H8, Ha, and textscii]]6584 or 1]6717,31)) commonly used in standard diagnostic diagrams."," The nature of our observations did not give us simultaneous access to the set of emission lines ]5007, $\beta$, $\alpha$, and ]6584 or ) commonly used in standard diagnostic diagrams."279" However, for all but two objects in our sample, the emission-line ratio N2=log([N1]6584/Ho) is lower than —0.5, with a median N2 value of —0.72."," However, for all but two objects in our sample, the emission-line ratio $\mathrm{N2} =280\log(\textrm\niia{}/\textrm\ha)$ is lower than $-0.5$, with a median $\mathrm{N2}$ value of $-0.72$."281 Such low values are indicative for a very low contamination by AGN in our sample (eg.?).., Such low values are indicative for a very low contamination by AGN in our sample \citep[eg.][]{bpt81}.282" For 24 galaxies of our sample we calculated, following ?,, the N2 ""concentrated ratio"", corresponding to the value in the nuclear region of the galaxy."," For 24 galaxies of our sample we calculated, following \citet{wright10}, the N2 “concentrated ratio”, corresponding to the value in the nuclear region of the galaxy."283" We defined the nuclear region as the spaxel with the highest flux along with its 8 nearest neighbours (corresponding to a 0.7"" diameter aperture, matching our mean spatial resolution of 0.65"")."," We defined the nuclear region as the spaxel with the highest flux along with its 8 nearest neighbours (corresponding to a $0.7''$ diameter aperture, matching our mean spatial resolution of $0.65''$ )."284" We assumed this aperture to be small enough to probe the inner nucleus part as objects usually span from 1"" to 2"" with SNR»2 in our observations.", We assumed this aperture to be small enough to probe the inner nucleus part as objects usually span from $1''$ to $2''$ with $SNR>2$ in our observations.285 Fig., Fig.286 2 shows i) the distribution of the 24 galaxies as a function of their global and nuclear N2 ratio (left panel) and ii) the relation between the global and nuclear N2 ratios for each galaxy (right panel)., \ref{distrib} shows i) the distribution of the 24 galaxies as a function of their global and nuclear N2 ratio (left panel) and ii) the relation between the global and nuclear N2 ratios for each galaxy (right panel).287 The median values of N2 for each distribution are not very different (A~ —0.07)., The median values of N2 for each distribution are not very different $\Delta\sim -0.07$ ).288 The median nuclear N2 ratio is lower than the global ratio which would not be the case if a significant fraction of our sample galaxies were hosting an AGN., The median nuclear N2 ratio is lower than the global ratio which would not be the case if a significant fraction of our sample galaxies were hosting an AGN.289" When comparing the global ratio to the nucleus ratio distribution, the highest bin does not shift and contains a single object."," When comparing the global ratio to the nucleus ratio distribution, the highest bin does not shift and contains a single object."290 We investigated in more detail the galaxy in this bin: VVDS140096645., We investigated in more detail the galaxy in this bin: VVDS140096645.291 It shows the following high N2 ratios: N2global=--0.252 and N2nuclear= —0.294., It shows the following high N2 ratios: $N2_\textrm{global}=-0.252$ and $N2_\textrm{nuclear}=-0.294$ .292" Looking further into its integrated spectrum (global and nuclear, see 3)), we noticed that: (i) the emission lines are broad, which is a possible sign of nuclear activity, (ii) the two nitrogen lines are clearly visible, as is the sulfur doublet,"," Looking further into its integrated spectrum (global and nuclear, see ), we noticed that: (i) the emission lines are broad, which is a possible sign of nuclear activity, (ii) the two nitrogen lines are clearly visible, as is the sulfur doublet,"293The binary content of a elobular cluster is important in determining the lrequency and nature of eluster stellar exotica. as well as the dvnamical evolution of the cluster.,"The binary content of a globular cluster is important in determining the frequency and nature of cluster stellar exotica, as well as the dynamical evolution of the cluster."294 ]t has long been recognized that binary formation is inevitable in a sell-gravitatingsvstem!., It has long been recognized that binary formation is inevitable in a self-gravitating.295. Indeed. (he presence of binaries as a central enerev source is vital to avoid complete (Goodman πι 1989).," Indeed, the presence of binaries as a central energy source is vital to avoid complete core-collapse (Goodman Hut 1989)."296 However. only more recently has it been realised that elobular clusters nist. also have formed with a sizeable binary population (see IIut et al.," However, only more recently has it been realised that globular clusters must also have formed with a sizeable binary population (see Hut et al."297 1992 [or an early review)., 1992 for an early review).298 That globular clusters harbour a mixture of dvnamically [ormed and primordial binaries can be used to understand observations of their stellar content. such as (he diverse blue straggler population in 47 Tucanae (Mapelli et al.," That globular clusters harbour a mixture of dynamically formed and primordial binaries can be used to understand observations of their stellar content, such as the diverse blue straggler population in 47 Tucanae (Mapelli et al."299 2004)., 2004).300 ]|xnowledge of the likely primordial binary fraction of globular clusters is essential as input to models of globular cluster evolution., Knowledge of the likely primordial binary fraction of globular clusters is essential as input to models of globular cluster evolution.301 H also provides a constraint on the cluster omnation process., It also provides a constraint on the cluster formation process.302 Considering (hat (he presence of binaries in (he cluster core has a pronounced effect on the core properties and cluster evolution (Hut 1996). knowledge of the central binary requency is also important.," Considering that the presence of binaries in the cluster core has a pronounced effect on the core properties and cluster evolution (Hut 1996), knowledge of the central binary frequency is also important."303 Indieations are (hat this is relatively small o£ the order of (e.g. Bellazzini et al., Indications are that this is relatively small – of the order of (e.g. Bellazzini et al.304 2002) or less (e.g. Cool Bolton 2002) — when compared to the yequencies of binaries observed in the solar neighbourhood (Duquennov Mayor 1991) and open clusters such as M67 (Fan et al., 2002) or less (e.g. Cool Bolton 2002) – when compared to the frequencies of binaries observed in the solar neighbourhood (Duquennoy Mayor 1991) and open clusters such as M67 (Fan et al.305 1996) which are of the order of50%., 1996) which are of the order of.306.. It would be particularly useful to take measurements of the current binary. traction in elobular clusters whether that be in (he core or outer regions and extrapolate backwards Lo gain a reliable determination of the primordial binary content., It would be particularly useful to take measurements of the current binary fraction in globular clusters – whether that be in the core or outer regions – and extrapolate backwards to gain a reliable determination of the primordial binary content.307 However. processes involved," However, processes involved"308show how the simultaneous comparison of the size and mass cistributions can reveal interesting insights on how to improve the performance of theoretical models of galaxy evolution.,show how the simultaneous comparison of the size and mass distributions can reveal interesting insights on how to improve the performance of theoretical models of galaxy evolution.309 We find. in agreement with previous studies. that this hierarchical nioclel provides a poor match to the size-mass relation of local galaxies. irrespective of the exact sample we compare it with.," We find, in agreement with previous studies, that this hierarchical model provides a poor match to the size-mass relation of local galaxies, irrespective of the exact sample we compare it with."310 In particular. the model tends to produce a much latter relation than the one actually observed.," In particular, the model tends to produce a much flatter relation than the one actually observed."311 This Hlattening is mainly. produced by the combined effects of having. with respect to the local data. too large (3 kpc) low-mass galaxies (<10: )). and of having a non-negligible fraction of compact. galaxies (0.51 kpc) at high masses (= LOMAL.)).," This flattening is mainly produced by the combined effects of having, with respect to the local data, too large $\sim 3$ kpc) low-mass galaxies $<10^{11}$ ), and of having a non-negligible fraction of compact galaxies $\lesssim 0.5-1$ kpc) at high masses $\gtrsim 10^{11}$ )."312 Such discrepancies are reflected. in. the predicted size distribution., Such discrepancies are reflected in the predicted size distribution.313 Although the model. produces a size clistribution in broad agreement with the data. it tends to overproduce the number of large galaxies bevond the peak £3 kpe). and the number of very compact. galaxies (S1 kpe).," Although the model produces a size distribution in broad agreement with the data, it tends to overproduce the number of large galaxies beyond the peak $\gtrsim 3$ kpc), and the number of very compact galaxies $\lesssim 1$ kpc)."314 We discussed that the former issue is present at all epochs. and it might therefore be linked to how spheroids ave formed in the first place. either. from not. properly reating initial disk instabilities and/or computing the sizes of remnants in gas-rich. mergers.," We discussed that the former issue is present at all epochs, and it might therefore be linked to how spheroids are formed in the first place, either from not properly treating initial disk instabilities and/or computing the sizes of remnants in gas-rich mergers."315 Regarding the overproduction of compact ancl massive (Mio(0.51)104 )) galaxies with respect to he data. already. pointed out in the recent Literature. we ind it to be less prominent than previously claimed. and confined to only ultracompact galaxies 2.0.5 kpe) when considering only ellipticals.," Regarding the overproduction of compact and massive $\sim (0.5-1)\times 10^{11}$ ) galaxies with respect to the data, already pointed out in the recent Literature, we find it to be less prominent than previously claimed, and confined to only ultracompact galaxies $\lesssim 0.5$ kpc) when considering only ellipticals."316 We discuss two possible reasons »hind the survival of such compact. galaxies until the esent epoch., We discuss two possible reasons behind the survival of such compact galaxies until the present epoch.317 First. we find that mocdel early-type galaxies end to be significantly. older than those in SDSS.," First, we find that model early-type galaxies tend to be significantly older than those in SDSS."318 This in urn mieht induce more compact galaxies at. fixed. stellar mass. given that galaxies formed at higher redshifts are more compact (see 2009a)).," This in turn might induce more compact galaxies at fixed stellar mass, given that galaxies formed at higher redshifts are more compact (see )."319 We also find that moclel earlv-tvpe compact galaxies underwent peculiar merging histories characterized by extremely compact. progenitors. that could. prevent them to elliciently. grow their sizes.," We also find that model early-type compact galaxies underwent peculiar merging histories characterized by extremely compact progenitors, that could prevent them to efficiently grow their sizes."320 ES acknowledges support from the Alexander von Humboldt Foundation anc partial support. from NASA Grant NNGOSGUTTG. MD is supported by NASA grant. LESA-NNGOGGCUIOG and NASA ADP/NNXOOADO2C. We thank CGuinevere Ixaulfmann. Ravi Sheth. Andrew Benson. Luigi Danese. Itosalind Skelton. Simon White. Qi Guo. and Volker Springel for various discussions.," FS acknowledges support from the Alexander von Humboldt Foundation and partial support from NASA Grant NNG05GH77G. MB is supported by NASA grant LTSA-NNG06GC19G and NASA ADP/NNX09AD02G. We thank Guinevere Kauffmann, Ravi Sheth, Andrew Benson, Luigi Danese, Rosalind Skelton, Simon White, Qi Guo, and Volker Springel for various discussions."321 We finally thank the referee [or several useful suggestions that improved the presentation of the paper., We finally thank the referee for several useful suggestions that improved the presentation of the paper.322 The filled squares in Figure Al represent the Dernardi οἱ al. (, The filled squares in Figure A1 represent the Bernardi et al. (3232009) estimate of the ERE obtained from the VYu method for our sample of carly-type galaxies selected: with concentration C.22.86.,2009) estimate of the ERF obtained from the $V/V_{\rm max}$ method for our sample of early-type galaxies selected with concentration $C_r>2.86$.324 For consistency. we here show that the V/V based ERP is consistent. within the errors. with the one obtained from the convolution of the luminosity function or velocity dispersion function with the bivariate distribution of points inthe £42. plane.," For consistency, we here show that the $V/V_{\rm max}$ -based ERF is consistent, within the errors, with the one obtained from the convolution of the luminosity function or velocity dispersion function with the bivariate distribution of points in the $L-R_e$ plane."325 More specifically. following the methods outlined in 7. and ?.. we have convolved the luminosity function ®(£) with the bivariate distribution of £; and I. llere £;; is the fraction of sources in the sample with effective radius A7. and luminosity Lj. normalized to the total number of sources with luminosity L;.," More specifically, following the methods outlined in \citet{Sheth03} and \citet{Shankar04}, we have convolved the luminosity function $\Phi(L)$ with the bivariate distribution of $L_j$ and $R_e$, Here $\xi_{ij}$ is the fraction of sources in the sample with effective radius $R_e^i$ and luminosity $L_j$, normalized to the total number of sources with luminosity $L_j$."326 The luminosity function (L) has been computed from the VYiuuu. method bv Bernardi ct al. (, The luminosity function $\Phi(L)$ has been computed from the $V/V_{\rm max}$ method by Bernardi et al. (3272009) for the same sample of galaxies. and we refer the reader to that paper for analvtical fits anc detailed discussions of the sample.,"2009) for the same sample of galaxies, and we refer the reader to that paper for analytical fits and detailed discussions of the sample."328 The result. of eq. CX1)), The result of Eq. \ref{eq|bivar}) )329 is shown in Figure X1 with long-dashed lines. which bracket the statistical uncertainties in the luminosity function Lit parameters.," is shown in Figure A1 with long-dashed lines, which bracket the statistical uncertainties in the luminosity function fit parameters."330 We have also used the bivariate distribution in equation CX1)) applied to the velocity dispersion. function Pia). again derived. by Bernardi et al. (," We have also used the bivariate distribution in equation \ref{eq|bivar}) ) applied to the velocity dispersion function $\Phi(\sigma)$, again derived by Bernardi et al. ("3312009). for. this same sample. and with the weights £;; now computed from the distribution of sources in the 0. &. plane.,"2009) for this same sample, and with the weights $\xi_{ij}$ now computed from the distribution of sources in the $\sigma-R_e$ plane."332 The result is shown with dotted lines in the same Figure. again bracketing the statistical uncertainties in the velocity dispersion function fit. parameters.," The result is shown with dotted lines in the same Figure, again bracketing the statistical uncertainties in the velocity dispersion function fit parameters."333 This exercise proves that. as expected. convolutions of other statistical distributions PCr) ," This exercise proves that, as expected, convolutions of other statistical distributions $\Phi(x)$ "334synchrotron self-compton (SSC) model (e.g.2).,synchrotron self-compton (SSC) model \citep[e.g.][]{maraschi92}.335 We assume a spherical. homogeneous emission region - coined blob - containing isotropically distributed non-thermal electrons and a randomly oriented magnetic field.," We assume a spherical, homogeneous emission region - coined blob - containing isotropically distributed non-thermal electrons and a randomly oriented magnetic field."336 DDue to the presence of this magnetic field the electrons emit svnchrotron radiation., Due to the presence of this magnetic field the electrons emit synchrotron radiation.337 The photons are then scattered off the same electron population via the inverse Compton process., The photons are then scattered off the same electron population via the inverse Compton process.338 The resulting spectrum shows the typical two bump structure commonly found in blazars., The resulting spectrum shows the typical two bump structure commonly found in blazars.339 In the following section the governing equations of the SSC-model are explained., In the following section the governing equations of the SSC-model are explained.340 To determine the time-dependent spectral energy distribution of blazars we solve the ditterential equation for the differential photon number density. obtained from the radiative transfer equation. including the corresponding terms with respect to SSC model. In the following context the well known ó-approximation (22). is applied to describe the synchrotron radiation in à convenient way.," To determine the time-dependent spectral energy distribution of blazars we solve the differential equation for the differential photon number density, obtained from the radiative transfer equation, including the corresponding terms with respect to SSC model, In the following context the well known $\delta$ -approximation \citep{fm66,341schlick02} is applied to describe the synchrotron radiation in a convenient way."342" Thus the synchrotron photon production rate Ay is given by with the pitch angle averaged total power P, emitted by a single electron having Lorentz factor y (2???) and y, being a function of v. obtained from the pitch angle averaged critical synchrotron frequency."," Thus the synchrotron photon production rate $R_S$ is given by with the pitch angle averaged total power $P_S$ emitted by a single electron having Lorentz factor $\gamma$ \citep{rybicki79, bg70, gs69}343 and $\gamma_c$ being a function of $\nu$, obtained from the pitch angle averaged critical synchrotron frequency."344 In optically thick regimes the emitted synchrotron radiation is absorbed by the emitting electrons itself., In optically thick regimes the emitted synchrotron radiation is absorbed by the emitting electrons itself.345 This is described by the synchrotron self absorption coethcient. which leads to the absorption rate The second main feature of the SSC model is Compton scattering of the synchrotron photons by the emitting electrons themselves.," This is described by the synchrotron self absorption coefficient, which leads to the absorption rate The second main feature of the SSC model is Compton scattering of the synchrotron photons by the emitting electrons themselves."346 Here the full Klein-Nishina cross section is used to calculate the photon production rate. The formula was taken from ? with minor corrections according to ?..," Here the full Klein-Nishina cross section is used to calculate the photon production rate, The formula was taken from \citet{wax05} with minor corrections according to \citet{cb90}."347 The photon energies are rewritten in terms of the electrons rest mass. so that /n*=enc? for the scattered photons and hv=eumc for the target photons.," The photon energies are rewritten in terms of the electrons rest mass, so that $h \nu = \epsilon m c^2$ for the scattered photons and $h \nu =348\epsilon_1 m c^2$ for the target photons."349" To make use of the full Klein- cross section we applied the approximate inverse Compton spectrum (2). of a single electron scattered off by a unit density photon field. where q""=eftdeyytl—e/yn and Αγ)<q.Ι."," To make use of the full Klein-Nishina cross section we applied the approximate inverse Compton spectrum \citep{jones68} of a single electron scattered off by a unit density photon field, where $ q''={\epsilon}/(4\epsilon_1 \gamma^2(1-\epsilon/\gamma))$ and $1/(4\gamma^2)<q''\leq 1$."350 This equation is valid Βου4eyy)., This equation is valid for.351". The corresponding ordinary Compton spectrum is approximately given by with -4y-e/e, and target photon energies in the range efdy. &e«e.", The corresponding ordinary Compton spectrum is approximately given by with $q'= 4 \gamma^2 \epsilon/\epsilon_1$ and target photon energies in the range $\epsilon_1/4\gamma^2\leq \epsilon < \epsilon_1$ .352 The last term describing the evolution of the photon number density represents the photons escape rate., The last term describing the evolution of the photon number density represents the photons escape rate.353 Here the photon escape time Έως 18 given by the light crossing time where A; is the radius of the emitting blob., Here the photon escape time $T_{esc}$ is given by the light crossing time where $R_b$ is the radius of the emitting blob.354 The escape time is chosen to be the light crossing time of the photons., The escape time is chosen to be the light crossing time of the photons.355" The time evolution of the electron distribution is described by the Kinetic equation The synchrotron loss is given by y,=Θες with the synchrotron power Pstv) (ef. 90100).", The time evolution of the electron distribution is described by the kinetic equation The synchrotron loss is given by $\dot \gamma_S=P_S(\gamma)/mc^2$ with the synchrotron power $P_S(\gamma)$ (cf. \ref{eq:PS}) )).356" Z4,==ΗΝο describes the electrons escaping from the emission region. where 7 isan empirical factor."," $T_{esc,e^-}= \eta R_b/c$ describes the electrons escaping from the emission region, where $\eta$ isan empirical factor."357 The inverse compton losses γιο including the full Klein-Nishina cross-section are adopted following 2.., The inverse compton losses $\dot \gamma_{IC}$ including the full Klein-Nishina cross-section are adopted following \citet{schlick02}. .358G. Clusters of galaxies provide a reservoir of barvons in the form of a hot plasma with typical temperatures of 10* 105 K. which emits over a broad baud frour Extreme Ultraviolet (EUV) to ~10 keV Xravs.,"G. Clusters of galaxies provide a reservoir of baryons in the form of a hot plasma with typical temperatures of $10^7$ $10^8$ K, which emits over a broad band from Extreme Ultraviolet (EUV) to $\sim 10$ keV X–rays."359 Most of the observed X-ray features of clusters can be well accounted for withiu the fraanework of thermal brensstrahluug enüssion plus chussion lines associated to the ietal coutent of the intracluster medi (CAD., Most of the observed X-ray features of clusters can be well accounted for within the framework of thermal bremsstrahlung emission plus emission lines associated to the metal content of the intra–cluster medium (ICM).360 However. a number of observations with the Extreme Ultraviolet Explorer (EUVE: e.g. Lieu et al," However, a number of observations with the Extreme Ultraviolet Explorer (EUVE; e.g. Lieu et al."361 1996a.b: Mittaz ct al.," 1996a,b; Mittaz et al."362 1998: Malonev Blaud-UWawthorn 2001). ROSAT (e.g. Donmaneute ct al.," 1998; Maloney Bland-Hawthorn 2001), ROSAT (e.g. Bonamente et al."363 2001a.b) and NMMNE-Newton (e.g. Finogucnoy et al.," 2001a,b) and XMM-Newton (e.g. Finoguenov et al."364 2003: IXaastra et al., 2003; Kaastra et al.365 2003) have claimed the detection of an excess of EUV aud soft N-vay ciuission in the spectrum of several clusters. with respect to what expected from a onetemperature plasima model.," 2003) have claimed the detection of an excess of EUV and soft X-ray emission in the spectrum of several clusters, with respect to what expected from a one–temperature plasma model."366" A uunber of sugeestions have been proposed for the origin of this OXCORS, the two nost popular scenarios bene the nonthermal origin from inverse Compton scattering of the cosmic microwave backeround photons by relativistic electrons in the intracluster eas (e.g. wane 1997: EnsBu Biermann 1998: De Paoli et al."," A number of suggestions have been proposed for the origin of this excess, the two most popular scenarios being the non–thermal origin from inverse Compton scattering of the cosmic microwave background photons by relativistic electrons in the intracluster gas (e.g. Hwang 1997; $\beta$ lin Biermann 1998; De Paolis et al."367" 2003). and the thermal origin due to warm gas at T~109A. either from diuside clisters or frou di""se filaments outside clusters (Liew et al."," 2003), and the thermal origin due to warm gas at $T\sim 10^6\,K$, either from inside clusters or from diffuse filaments outside clusters (Lieu et al."368 1996b: Nevalainen et al., 1996b; Nevalainen et al.369 2003)., 2003).370 However. the existence of the excess js been disputed by Bowver ct al. (," However, the existence of the excess has been disputed by Bowyer et al. ("3711999). who argued hat the EUV excess is an artifact caused bv improper subtraction of the instrumental background (see also Berghoter Bowyer 2002: Duxret et al.,"1999), who argued that the EUV excess is an artifact caused by improper subtraction of the instrumental background (see also Berghofer Bowyer 2002; Durret et al."372 2002). although they conclided that a relatively weak EUV excess in the Virgo and Coma clusters may be real.," 2002), although they concluded that a relatively weak EUV excess in the Virgo and Coma clusters may be real."373 Bregman ct al. (, Bregman et al. (3742001) also argued that the excess may be caused by an improper inclusion in the data analysis of the effect of fluctuations of the ealactic hydrogen column deusity.,2004) also argued that the excess may be caused by an improper inclusion in the data analysis of the effect of fluctuations of the galactic hydrogen column density.375 Even within the framework of the thermal models. it is a matter of debate whether the gas responsible for the excess ds located within clusters or in largescale filamentary structures.," Even within the framework of the thermal models, it is a matter of debate whether the gas responsible for the excess is located within clusters or in large–scale filamentary structures."376 For ustauce. Ettori (2003) found evidence for as much as LF per cout of the barvons iu Justers to be prestumably in the form of win (107 LO’ IS) naterial which must enüt EUV or soft Nouv photons in excess to those expected frou the hot phase of the ICAL," For instance, Ettori (2003) found evidence for as much as $17$ per cent of the baryons in clusters to be presumably in the form of warm $10^5$ $10^7$ K) material, which must emit EUV or soft X-ray photons in excess to those expected from the hot phase of the ICM."377 On the other hand. Naastra et al. (," On the other hand, Kaastra et al. ("3782003) and Finoguenov ct al. (,2003) and Finoguenov et al. (3792003) claim that the soft excess may have originated in flaments in the eas distribution im the vicinity of clusters.,2003) claim that the soft excess may have originated in filaments in the gas distribution in the vicinity of clusters.380 Tu this paper. we present an analysis of a set of clusters extracted from a large cosmological hivdrodyvuaiical «αμαπο (Boregani et al.," In this paper, we present an analysis of a set of clusters extracted from a large cosmological hydrodynamical simulation (Borgani et al."381 2001. Paper D. which is aimed at investigating the presence aud origi of a soft Xray OXCOSS in them spectra.," 2004, Paper I), which is aimed at investigating the presence and origin of a soft X--ray excess in their spectra."382 By taking advantage of the cosmological environment of our sinulatiou. which includes radiative cooling. star formation and ealactic winds trigeered. by supernova (SN). we “observe” clusters in projection and estimate their spectra also includiug the contribution from the backgrouud/foreeround largescale gas distribution.," By taking advantage of the cosmological environment of our simulation, which includes radiative cooling, star formation and galactic winds triggered by supernova (SN), we “observe” clusters in projection and estimate their spectra also including the contribution from the background/foreground large--scale gas distribution."383 Since iu our simulation we treat oulv thermal cuussivity processes. the two main questions that we intend to address are the following: (a) does a realistic description of the evolution of cosmic barvous account for a soft Nrav excess of thermal origin as large as that observed in the spectra of clusters? (," Since in our simulation we treat only thermal emissivity processes, the two main questions that we intend to address are the following: (a) does a realistic description of the evolution of cosmic baryons account for a soft X–ray excess of thermal origin as large as that observed in the spectra of clusters? ("384b) is the excess associated to wari eas residing within clusters or to flaments observed in projection?,b) is the excess associated to warm gas residing within clusters or to large--scale filaments observed in projection?385"The curves in Figure 6 show these two quantities at several different ionized fractions (logzx;= —4,—3,—2, and —1, from top to bottom in the left hand panel).","The curves in Figure \ref{fig:int-spec} show these two quantities at several different ionized fractions $\log x_i=-4,\,-3,\,-2,$ and $-1$, from top to bottom in the left hand panel)."386 The solid squares show the asymptotic high-energy estimates from the ? fitting formulae. (, The solid squares show the asymptotic high-energy estimates from the \citet{shull85} fitting formulae. (387These are the terms in eqs. 13-,These are the terms in eqs. \ref{eq:gnedin-fit1}-388-14 without any energy dependence.), \ref{eq:gnedin-fit2} without any energy dependence.)389" These fitting formulae provide reasonable order-of-magnitude estimates so long as Emin=300eV (or slightly higher if z;~ 0.1), although they systematically underestimate fLy and fios.n."," These fitting formulae provide reasonable order-of-magnitude estimates so long as $E_{\rm min} \ga 300 \eV$ (or slightly higher if $x_i \sim 0.1$ ), although they systematically underestimate $\bar{f}_{\rm Ly\alpha}$ and $\bar{f}_{\rm ion,HI}$."390" These deviations are most significant at small ionized fractions, where heating is least important."," These deviations are most significant at small ionized fractions, where heating is least important."391" We therefore again recommend interpolation of the exact results for high-accuracy work, especially if the heating by soft X-rays E<300eV is included."," We therefore again recommend interpolation of the exact results for high-accuracy work, especially if the heating by soft X-rays $E \la 300 \eV$ is included."392" Using a Monte Carlo model, we have re-examined the fate of fast electrons scattering through a background gas of primordial origin."," Using a Monte Carlo model, we have re-examined the fate of fast electrons scattering through a background gas of primordial origin."393" We included electron-electron scattering as well as collisional ionization and excitation of HI, Hel, and Hell, explicitly tracking all levels up to n.=4 and using an analytic extrapolation to higher levels."," We included electron-electron scattering as well as collisional ionization and excitation of HI, HeI, and HeII, explicitly tracking all levels up to $n=4$ and using an analytic extrapolation to higher levels."394" We separately followed all excitations producing HI photons, which can be important in modeling the observable properties of the IGM at high redshifts (see ? and ?,, for example) and have not been explicitly tracked previously except at the highest energies."," We separately followed all excitations producing HI photons, which can be important in modeling the observable properties of the IGM at high redshifts (see \citealt{kuhlen06-21cm} and \citealt{furl06-review}, for example) and have not been explicitly tracked previously except at the highest energies."395 We used recent calculations of ionization and excitation cross-sections at /;<1keV and extrapolated to higher energies using the Bethe approximation., We used recent calculations of ionization and excitation cross-sections at $E<1 \keV$ and extrapolated to higher energies using the Bethe approximation.396" In highly neutral gas (7; 10” we found that ~20% of the electron energy is deposited as heat,5), with the remainder split roughly equally between ionization and excitation."," In highly neutral gas $x_i \la 10^{-3}$ ), we found that $\sim 20\%$ of the electron energy is deposited as heat, with the remainder split roughly equally between ionization and excitation."397" In this regime, the results are not strongly sensitive to ση, at least at high energies, because most of the heating comes from secondary electrons, with energies below 10eV."," In this regime, the results are not strongly sensitive to $x_i$, at least at high energies, because most of the heating comes from secondary electrons, with energies below $10 \eV$."398" At higher x;, the heating fraction rises rapidly, exceeding ~65% by x;~0.1."," At higher $x_i$, the heating fraction rises rapidly, exceeding $\sim 65\%$ by $x_i \sim 0.1$."399" We find that the excitation and ionization energy deposition rates are always comparable, and that ~80% of the excitation energy goes into HI regardless of electron energy and z;."," We find that the excitation and ionization energy deposition rates are always comparable, and that $\sim 80\%$ of the excitation energy goes into HI regardless of electron energy and $x_i$."400" Although our calculations used parameters appropriate to the low-density IGM, our results may also be applied to denser systems, because the density only enters through the Coulomb logarithm affecting the electron-electron scattering rate."," Although our calculations used parameters appropriate to the low-density IGM, our results may also be applied to denser systems, because the density only enters through the Coulomb logarithm affecting the electron-electron scattering rate."401 Varying the density of target atoms by many orders of magnitude only affects the energy deposition fractions by a few percent in absolute terms., Varying the density of target atoms by many orders of magnitude only affects the energy deposition fractions by a few percent in absolute terms.402" We also note that, when collective plasma effects are included, the background temperature becomes irrelevant for the electron energies under consideration (??).."," We also note that, when collective plasma effects are included, the background temperature becomes irrelevant for the electron energies under consideration \citep{schunk71,xu91}."403" For the most part, our results agree with previous estimates from ? and ?, although we have found some discrepancies with the commonly-used fitting formulae from the former."," For the most part, our results agree with previous estimates from \citet{shull85} and \citet{xu91}, although we have found some discrepancies with the commonly-used fitting formulae from the former."404" In general, at high energies their results slightly underestimate the importance of collisional excitation but are otherwise accurate."," In general, at high energies their results slightly underestimate the importance of collisional excitation but are otherwise accurate."405" At lower energies, the differences in cross-sections become more important and the discrepancies increase (at least with reference to the fitting formulae of ?))."," At lower energies, the differences in cross-sections become more important and the discrepancies increase (at least with reference to the fitting formulae of \citealt{ricotti02}) )."406" However, our results show substantial differences with those of ?,, who examined the high-energy limit with a Monte Carlo model similar to ours."," However, our results show substantial differences with those of \citet{valdes08}, who examined the high-energy limit with a Monte Carlo model similar to ours."407" These discrepancies are especially large at moderate and high ionized fractions, where we find substantially more heating and less ionization and excitation."," These discrepancies are especially large at moderate and high ionized fractions, where we find substantially more heating and less ionization and excitation."408 We also find that a higher fraction of excitation energy is deposited in HI photons., We also find that a higher fraction of excitation energy is deposited in HI photons.409" The latter is probably due to our different treatments of collisional excitation (in particular, their neglect of excitations to states other than the np sublevels), but the source of the former is unclear."," The latter is probably due to our different treatments of collisional excitation (in particular, their neglect of excitations to states other than the $np$ sublevels), but the source of the former is unclear."410" In any case, we advocate interpolation of the exact results when high accuracy is necessary, especially because the energy dependence is quite significant."," In any case, we advocate interpolation of the exact results when high accuracy is necessary, especially because the energy dependence is quite significant."411" Our detailed numerical results are available upon request, including tables for the energy deposition"," Our detailed numerical results are available upon request, including tables for the energy deposition"412system aud he role played by blacx hole spin in the enerey of the jets.,system and the role played by black hole spin in the energy of the jets.413 Given tie ability to stunlate he relativistic dyuaimics of the inner engine of a collapsar. we see& lo answer three questions: what is the range o‘Lorentz [actors in the jets: how much energy is trausportect N these jets (aux for |ow loug): :uxl how variable is the flux of energy.," Given the ability to simulate the relativistic dynamics of the inner engine of a collapsar, we seek to answer three questions: what is the range of Lorentz factors in the jets; how much energy is transported by these jets (and for how long); and how variable is the flux of energy."414" The answer to these questious should help to estallisi whetLer he dynamics of the GRMHD interaction between αι»ς, black |ole. aud jet in ow numerical setup support the basic tenets of the collapsar moclel."," The answer to these questions should help to establish whether the dynamics of the GRMHD interaction between disk, black hole, and jet in our numerical setup support the basic tenets of the collapsar model."415 Since the GRALHD cocle is scae-[ree. we will cliscuss the scaling of our results to the collapsar model in a separate section.," Since the GRMHD code is scale-free, we will discuss the scaling of our results to the collapsar model in a separate section."416 We follow the notation usec in DHO3: a stummary of the equations evolved by the GRMHD code atd tlie set of dynamical variales available for analysis is given in the appendix. where a cliscusSLOLL O ‘the numerical determination of the Lorentz factor is also giveu.," We follow the notation used in DH03; a summary of the equations evolved by the GRMHD code and the set of dynamical variables available for analysis is given in the appendix, where a discussion of the numerical determination of the Lorentz factor is also given."417 We repor all sitjulation results in geometrodynamic units (Misuer. Thorne Wheeler. 1973). so that time «vnc cisance are measured in units of the black hole mass. A/pyy.," We report all simulation results in geometrodynamic units (Misner, Thorne Wheeler, 1973), so that time and distance are measured in units of the black hole mass, $M_{\rm BH}$."418 The GRMHD code uses the test-fuid ap»roximation. so that euergy of the orbiing fluid does not alter the background spacetline.," The GRMHD code uses the test-fluid approximation, so that energy of the orbiting fluid does not alter the background spacetime."419 Wie express Inass and energy in ‘elation to the maximum initial torus values., We express mass and energy in relation to the maximum initial torus values.420 The simulation ouput cosists of an exteusive collection of data ciups. obtained at increments of 2Al of simulation ime: {1ese clips are used to COLipute shell-averaged fluxes to measure euergy output aud teuporal variability (see DHIN for details).," The simulation output consists of an extensive collection of data dumps, obtained at increments of $2\,M$ of simulation time; these dumps are used to compute shell-averaged fluxes to measure energy output and temporal variability (see DHK for details)."421 The initial state of our simllatious is similar to that of DHIx: a torus with a uear-Ixepleriau 'otation profile is seeded with a weak poloidal magnetic field (the streneth of the field is set using he ratio of average gas to maegletic pressure in the torus. Aq=((Lomas2VÉmag) 100) to trigger he MRI aid generate au accretion flow.," The initial state of our simulations is similar to that of DHK: a torus with a near-Keplerian rotation profile is seeded with a weak poloidal magnetic field (the strength of the field is set using the ratio of average gas to magnetic pressure in the torus, $\beta_{\rm disk}=\langle P_{\rm gas}\rangle/\langle P_{\rm mag}\rangle 422=$ 100) to trigger the MRI and generate an accretion flow."423 As noted above. two important code improvements allow is to probe the dynamics of the axial funnel more thoroughly: an expanded dyuamic range aud a significantly iucreased ceiling oi the Lorentz factor (uow at 20).," As noted above, two important code improvements allow us to probe the dynamics of the axial funnel more thoroughly: an expanded dynamic range and a significantly increased ceiling on the Lorentz factor (now at 50)."424 These improvements are crucial o addressing issues of jet dyvuanites., These improvements are crucial to addressing issues of jet dynamics.425 Firhermore. to follow the evolution of the jets away from the inuer engine. we use a greatly extedec radial range (in the 2D simulations).," Furthermore, to follow the evolution of the jets away from the inner engine, we use a greatly extended radial range (in the 2D simulations)."426 Also. the grid outside he initial torus contains a racially infalliig cust (the Boudi solutiou discussed in Hawley. Smarr Wilson. 1981: hereafter HSW) o provide au influx of matter against which the eimergiug jets nust compete.," Also, the grid outside the initial torus contains a radially infalling dust (the Bondi solution discussed in Hawley, Smarr Wilson, 1984; hereafter HSW) to provide an influx of matter against which the emerging jets must compete."427 The deusity contrast between dust and torus. at 10.5 (see Table 1)). was chosen to je larger than the typical value of a collapsar environment (W93) for reasons that will be touched pon i the following sections.," The density contrast between dust and torus, at $10^{-6}$ (see Table \ref{params}) ), was chosen to be larger than the typical value of a collapsar environment (W93) for reasons that will be touched upon in the following sections."428 Iu some simulations. this torus is also embedded. in au external vertica| maguetic field (Wald. 197E: hereafter WT1) to gain insights into the possible effects of such ainbieit fields.," In some simulations, this torus is also embedded in an external vertical magnetic field (Wald, 1974; hereafter W74) to gain insights into the possible effects of such ambient fields."429 The streneth of these fields was chosen so that they could potentially interact witli the MBI-geuerated fields. as measure by μμ. he ratio of gas to magnetic pressure outside the initial orus.," The strength of these fields was chosen so that they could potentially interact with the MRI-generated fields, as measured by $\beta_{\rm dust}$, the ratio of gas to magnetic pressure outside the initial torus."430" The choice of initial streneth (see Table 1)) was established in tests which revealed tliat strong fields teudecl to disrupt the initial torus through maguetic teusiou: ""óiutermediate"" values wele ¢10561 for these simulatious.", The choice of initial strength (see Table \ref{params}) ) was established in tests which revealed that strong fields tended to disrupt the initial torus through magnetic tension; “intermediate” values were chosen for these simulations.431solution. the proper energy density. Lorentz factor and proper number density of the shocked fluid are given by where [ is the Lorentz factor of the shock. and The 4 coordinate of a fluid element is given by where Ro and fo are the shock radius and coordinate time. respectively. when the fluid element crosses the shock.,"solution, the proper energy density, Lorentz factor and proper number density of the shocked fluid are given by where $\Gamma$ is the Lorentz factor of the shock, and The $\chi$ coordinate of a fluid element is given by where $R_0$ and $t_0$ are the shock radius and coordinate time, respectively, when the fluid element crosses the shock."432" Since DI?x(7. we obtain that Using equation 19 and the relation dr!=dr/>. we can write equation 13. interms of 4: Solving equation 20 we obtain where 5,0—σαν=1) is the initial Lorentz factor of the electron. just the shock. and u 18 the maximal Lorentz factor at behindV>1. which corresponds to an electron with σος>x. and ts given by The of electrons with a Lorentz factor within the interval fraction[55.54alld5,] is given by: withN(z,)d5,/n. and remains constant as these quantities evolve increasing 4."," Since $\Gamma^2\propto t^{k-3}$, we obtain that Using equation \ref{to_chi} and the relation $dt'=dt/\gamma$, we can write equation \ref{dgamma_dt} interms of $\chi$: Solving equation \ref{dgamma_dchi} we obtain where $\gamma_{e,0}\equiv\gamma_e(\chi=1)$ is the initial Lorentz factor of the electron, just behind the shock, and $\gamma_{\rm max}(\chi)$ is the maximal Lorentz factor at $\chi>1$, which corresponds to an electron with $\gamma_{e,0}\to \infty$, and is given by The fraction of electrons with a Lorentz factor within the interval $[\gamma_{e},\gamma_{e}+ d\gamma_{e}]$ is given by: $N(\gamma_{e})d\gamma_{e}/n$, and remains constant as all these quantities evolve with increasing $\chi$."433 The electron distribution ts therefore given by: Where5uin(A2=Femina.\)-," The electron distribution is therefore given by: where $\gamma_{\rm min}(\chi)=\gamma_e(\gamma_{{\rm min},0},\chi)$."434 We now have explicit expressions for both. the hydrodynamical quantities and the electron distribution. over all relevant space-time. and can calculate the flux density near the various break frequencies.," We now have explicit expressions for both the hydrodynamical quantities and the electron distribution, over all relevant space-time, and can calculate the flux density near the various break frequencies."435" For breaks that are in the opticallythin regime (b=2.3.9.11) one may use the equation which is a generalization of equation 13 of GPS99a. where d; and z are the Juminosity distance and cosmological redshift of the source. respectively. P7, is the radiated power per unit volume per unit frequency in the local rest frame of the fluid. and should be taken at the coordinate time ¢=f-ric. where t-=fasfCl+2). E is the energy of the blast wave. v=R/R, (e.g. GPS99a). i!=vr(l—jr)."," For breaks that are in the opticallythin regime (b=2,3,9,11) one may use the equation which is a generalization of equation 13 of GPS99a, where $d_L$ and $z$ are the luminosity distance and cosmological redshift of the source, respectively, $P'_{\nu'}$ is the radiated power per unit volume per unit frequency in the local rest frame of the fluid, and should be taken at the coordinate time $t=t_z+r\mu/c$, where $t_z\equiv t_{\rm obs}/(1+z)$, $E$ is the energy of the blast wave, $y\equiv R/R_l$ (e.g. GPS99a), $\nu'=\nu\gamma(1-\beta\mu)$."436" The Spectral emissivity of a single electron (in the fluid rest frame) ts given by where g, is the electric charge of the electron. à is the pitch angle between the direction of the electron’s velocity and the magnetic field. in the local rest frame of the fluid. and F is the standard synchrotron function (e.g. Rybicki Lightman 1979)."," The Spectral emissivity of a single electron (in the fluid rest frame) is given by where $q_e$ is the electric charge of the electron, $\alpha$ is the pitch angle between the direction of the electron's velocity and the magnetic field, in the local rest frame of the fluid, and $F$ is the standard synchrotron function (e.g. Rybicki Lightman 1979)."437" In order to obtain an expression for P7, (which appears in equation 24)) we average P7,. over o. assuming an isotropic distribution of electrons in the local rest frame. and then integrate over the electron distribution. For the remaining spectral breaks (bz1.4.5.6.7.8.10). where the system Is not always optically thin. we follow the formalism of GPS99b."," In order to obtain an expression for $P'_{\nu'}$ (which appears in equation \ref{optically_thin}) ) we average $P'_{\nu',e}$ over $\alpha$, assuming an isotropic distribution of electrons in the local rest frame, and then integrate over the electron distribution, For the remaining spectral breaks (b=1,4,5,6,7,8,10), where the system is not always optically thin, we follow the formalism of GPS99b."438" Since the emission is isotropic in the local rest frame of the fluid. the emission coefficient is simply Επ. where P7, is given by equation 29.."," Since the emission is isotropic in the local rest frame of the fluid, the emission coefficient is simply $j'_{\nu'}=P'_{\nu'}/4\pi$ , where $P'_{\nu'}$ is given by equation \ref{Pnu}. ."439 The absorption coefficient is given by Since the flow ts spherically symmetric. the afterglow image is circular. with physical radius of," The absorption coefficient is given by Since the flow is spherically symmetric, the afterglow image is circular, with physical radius of"440"The recent abundance of cosinological data iu the last few decades has provided compelling evidence towards a standard concordance cosmology. in which the Universe is composed of approxtmately barvons. ‘dark’ matter and ""dark energy.","The recent abundance of cosmological data in the last few decades has provided compelling evidence towards a standard concordance cosmology, in which the Universe is composed of approximately baryons, `dark' matter and `dark' energy."441 One of the main challenges of modera cosmology is to understand the nature of the iivsterious dark energy which drives the observed cosmic acceleration , One of the main challenges of modern cosmology is to understand the nature of the mysterious dark energy which drives the observed cosmic acceleration .442The Inteerated Sacls-Wolte CSW) effect ds a secondary anidsotropv οἳ the Cosmic Microwave Backeround (CAIB). which arises because of he variation with time of the cosmuc gravitational potential between ocal observers and the surface of last scattering.," The Integrated Sachs-Wolfe (ISW) effect is a secondary anisotropy of the Cosmic Microwave Background (CMB), which arises because of the variation with time of the cosmic gravitational potential between local observers and the surface of last scattering."443 The oteutial can be traced by Large Scale Structure (LSS) survevs (7).. and the ISW effect is therefore a probe which inks he high redshift CAIB with the low redshift matter distribution aud cau be detected by cross-correlating the Wo.," The potential can be traced by Large Scale Structure (LSS) surveys , and the ISW effect is therefore a probe which links the high redshift CMB with the low redshift matter distribution and can be detected by cross-correlating the two."444" As a cosinological probe. the ISW effect has less statistical power than weak lensing or galaxy clustering 7). but it is directly sensitive to dark enerev. curvature or modified eravity (2?777).. such tha in universes where io0dified gravity and curvature excluded. detection of the ISW signal provides a direc sjeuature of dark οποίον,"," As a cosmological probe, the ISW effect has less statistical power than weak lensing or galaxy clustering , but it is directly sensitive to dark energy, curvature or modified gravity , such that in universes where modified gravity and curvature are excluded, detection of the ISW signal provides a direct signature of dark energy."445 Ta more geucra uuverses. the ISW effect can be used to trace alternative models of eyavity.," In more general universes, the ISW effect can be used to trace alternative models of gravity."446 The CAMB WMAP survey is already optimal for detecting the ISW signa (sce Sections 77. aud ?7)). ane sienificance is not expected to ierease with the arriva of Planck. unless the effect of the foreground Galactic mask can be reduced.," The CMB WMAP survey is already optimal for detecting the ISW signal (see Sections \ref{sec:theory} and \ref{sec:method}) ), and significance is not expected to increase with the arrival of Planck, unless the effect of the foreground Galactic mask can be reduced."447 The amplitude of the measured ISW signal should however depend strongly on the details of the local tracer of mass., The amplitude of the measured ISW signal should however depend strongly on the details of the local tracer of mass.448 Survey optuuisations show that an ideal ISW survey requires the same configuration as surveys which are optimised for weak lensing or galaxy clustering - caning that an optimal measure of the ISW signal will essentially come “for free’ with future planned weak lensing aud galaxy cblusterime surveys?)., Survey optimisations show that an ideal ISW survey requires the same configuration as surveys which are optimised for weak lensing or galaxy clustering - meaning that an optimal measure of the ISW signal will essentially come `for free' with future planned weak lensing and galaxy clustering surveys.449. In the best scenario. a lo detection is expected(?).. aud it has been shown that combined with weak lensing. galaxy correlation and other probes such as clusters. the ISW can be useful to break parameter degeneracies(?).. making it a promising probe.," In the best scenario, a $\sigma$ detection is expected, and it has been shown that combined with weak lensing, galaxy correlation and other probes such as clusters, the ISW can be useful to break parameter degeneracies, making it a promising probe."450 Iuitial attempts to detect the ISW effect with COBE as the CAIB tracer were fruitless(2).. but since the arrival of WALAP data. tens of positive detections lave been made. with the highest significance reported for analyses usimg a tomographic combination of surveys (soe Sections 77 and 2? for a cletailed review of detections).," Initial attempts to detect the ISW effect with COBE as the CMB tracer were fruitless, but since the arrival of WMAP data, tens of positive detections have been made, with the highest significance reported for analyses using a tomographic combination of surveys (see Sections \ref{sec:theory} and \ref{sec:method} for a detailed review of detections)."451 However. severa studies using the same tracer of LSS appear to have contradicting conclusions. some analvses do not find correlation where others do. aud as statistical methods to analyse the data evolve. the significance of the ISW signal Is sometimes reduced222).," However, several studies using the same tracer of LSS appear to have contradicting conclusions, some analyses do not find correlation where others do, and as statistical methods to analyse the data evolve, the significance of the ISW signal is sometimes reduced."452 Iu Section ?7.. we describe the cause of the ISW effect and review curent detectious.," In Section \ref{sec:theory}, we describe the cause of the ISW effect and review current detections."453" Iu Section οον, we describe the methodology for detection and measuring the ISW signal and review a laree proportion of reported detections 1i the literature. as well as their advantages and disadvantages."," In Section \ref{sec:method}, we describe the methodology for detection and measuring the ISW signal, and review a large proportion of reported detections in the literature, as well as their advantages and disadvantages."454 Tavine ideutified the main issues with current methods. we propose a new aud complete method in Section ??.. which capitalises ou the fact that different st:istical methods are complementary aud uses sparse inpaiutiug to solve the issue of missing data and a bootstrapping technique tomeasure the estimuitor's probability distribution function (PDF).," Having identified the main issues with current methods, we propose a new and complete method in Section \ref{sec:saclaymethod}, which capitalises on the fact that different statistical methods are complementary and uses sparse inpainting to solve the issue of missing data and a bootstrapping technique tomeasure the estimator's probability distribution function (PDF)."455 Iu Section 5.. we validate our new icthod using simulations for 2\TASS and Enclid-like surveys.," In Section \ref{sec:validation}, we validate our new method using simulations for 2MASS and Euclid-like surveys."456 In Section 6.. we apply our new method to WALAP 7 and the 2\LASS survey.," In Section \ref{sec:data}, we apply our new method to WMAP 7 and the 2MASS survey."457 In Section 77S we present our conclusions.," In Section \ref{sec:discussion}, , we present our conclusions."458"an ""outside-in"" twpe comes from examining the individual Πειραιά anc V-band outburst light. curves in bie.",an “outside-in” type comes from examining the individual R-band and V-band outburst light curves in Fig.459 ὸ and Fig., \ref{fig:rband} and Fig.460 9 respectively., \ref{fig:vband} respectively.461 Lt can easily be seen that the rise to outburst eclipses are wider than the eclipses in the late stages of the outburst (sce Fig., It can easily be seen that the rise to outburst eclipses are wider than the eclipses in the late stages of the outburst (see Fig.462 2. Fig 3 and Fig.," 2, Fig 3 and Fig."463 6).Xlso. eclipse depths during the rise are much shallower than the eclipses during the decline. indicating that during mid eclipse there is more (ux observed. from either side of the secondary star during the rise phase than the decline phase of the outburst.," 6).Also, eclipse depths during the rise are much shallower than the eclipses during the decline, indicating that during mid eclipse there is more flux observed from either side of the secondary star during the rise phase than the decline phase of the outburst."464 Since the outburst eclipses are extremely symmetric. we fitted them with a Gaussian to measure the curation and depth of cach eclipse.," Since the outburst eclipses are extremely symmetric, we fitted them with a Gaussian to measure the duration and depth of each eclipse."465 The residuals from the Caussian [it were at the level., The residuals from the Gaussian fit were at the level.466 Fable 1. lists the full width half minimum for each light curve along with the depth of each eclipse., Table \ref{tab:observation_log} lists the full width half minimum for each light curve along with the depth of each eclipse.467 Lt is evident from Table 1 that the outburst eclipses increase their depth anc become narrower as the eruption progresses., It is evident from Table \ref{tab:observation_log} that the outburst eclipses increase their depth and become narrower as the eruption progresses.468 I is a further indication that the outburst starts at the edge of the disc. whose luminous part is relatively large at the beginning of the outburst. and decreases in size as the outburst progresses.," It is a further indication that the outburst starts at the edge of the disc, whose luminous part is relatively large at the beginning of the outburst, and decreases in size as the outburst progresses."469 Patterson(1981) observed ανα nova oscillations (DNO) during an eruption of WP Cas in LOTS., \scite{Patterson81} observed dwarf nova oscillations (DNO) during an eruption of HT Cas in 1978.470 X. power spectrum of our V. and. 1t out-of-eclipse data did. not reveal any significant oscillations., A power spectrum of our V and R out-of-eclipse data did not reveal any significant oscillations.471 Given that we have a relatively low temporal resolution and. also that DNOs are usually found in shorter wavelength regions. this is not very surprising.," Given that we have a relatively low temporal resolution and also that DNOs are usually found in shorter wavelength regions, this is not very surprising."472 ‘Table 1. shows the mid-eclipse timines obtained during the quiescent ancl outburst. periods in 1995 and. 1997., Table \ref{tab:observation_log} shows the mid-eclipse timings obtained during the quiescent and outburst periods in 1995 and 1997.473 Previous data. that were explicitly tabulated. from Patterson(1981).. Zhang.Itobinson&Nather(1986).. and Horne (1991).," Previous data, that were explicitly tabulated, from \scite{Patterson81}, \scite{Zhang86}, and Horne (1991)."474 were converted from LHeliocentrie Julian Date to Barycentric Dynamical Julian Date ancl used in our calculations., were converted from Heliocentric Julian Date to Barycentric Dynamical Julian Date and used in our calculations.475 Only the quiescent data were used in order to caleulate the refined ephemeris presented. here., Only the quiescent data were used in order to calculate the refined ephemeris presented here.476 “Phe mid-eclipse positions were calculated by measuring the time of the white dwarl ingress and egress., The mid-eclipse positions were calculated by measuring the time of the white dwarf ingress and egress.477 For the outburst data presented. here we used. a Gaussian curve fit to measure the time of minimum Lux., For the outburst data presented here we used a Gaussian curve fit to measure the time of minimum flux.478 Vhe data for November 17th from Ixeele Observatory have a timing problem and they have been left out from all our calculations involving time and are not displaved on the O-C cliagram., The data for November 17th from Keele Observatory have a timing problem and they have been left out from all our calculations involving time and are not displayed on the O-C diagram.479 3elow ds our calculated. ephemeris in. Baryeentric Dynamical Julian Date with the uncertainties shown in brackets., Below is our calculated ephemeris in Barycentric Dynamical Julian Date with the uncertainties shown in brackets.480 We used our derived. period to project back to the Zi quoted by Patterson(1981)., We used our derived period to project back to the $T_{0}$ quoted by \scite{Patterson81}.481. Phe residuals of our ephemeris calculations (O-C€ diagram) can be seen in Fig. 4.., The residuals of our ephemeris calculations (O-C diagram) can be seen in Fig. \ref{fig:o-c}.482 The unused: outburst eclipse points are represented: as circles., The unused outburst eclipse points are represented as circles.483 Woodetal.(1995). speculated that IE. Cas might have a non zero period derivative., \scite{Wood95} speculated that HT Cas might have a non zero period derivative.484 Llowever. due to the large scatter of the data points in Fig.," However, due to the large scatter of the data points in Fig."485 4 we could only fit the data with a linear ephemeris., \ref{fig:o-c} we could only fit the data with a linear ephemeris.486 During the outburst an anomaly of the O-C points is observed. especially during the second night of observations when the svstem is at its carly decline stages. (," During the outburst an anomaly of the O-C points is observed, especially during the second night of observations when the system is at its early decline stages. ("487see Fig. 5)),see Fig. \ref{fig:outburst_o-c}) )488 During the decline from outburst the points gradually cirift o à maximum cillerence of TO seconds from the caleulated value., During the decline from outburst the points gradually drift to a maximum difference of 79 seconds from the calculated value.489 Observations of other objects on the same night show his is not à computer clock problem., Observations of other objects on the same night show this is not a computer clock problem.490 The O-C points then return to their normal level the following night., The O-C points then return to their normal level the following night.491 We comment on this unusual behaviour in the discussion section of the )Aper., We comment on this unusual behaviour in the discussion section of the paper.492 A model that takes into account the [Lux emitted by a white dwarl encirclec by an optically thick accretion disc and a Roche-lobe filling secondary star. which is. tically locked in the frame of the binary. was mocified ancl usec for the investigation of the observed. light curves during the ITE Cas outburst.," A model that takes into account the flux emitted by a white dwarf encircled by an optically thick accretion disc and a Roche-lobe filling secondary star, which is tidally locked in the frame of the binary, was modified and used for the investigation of the observed light curves during the HT Cas outburst."493 The model also accounts for the irradiation of the secondarys surface by a heating source at the centre of the disc., The model also accounts for the irradiation of the secondary's surface by a heating source at the centre of the disc.494 We also incorporate approximate temperature dependent limb darkening coellicients taken from Al-Naimiv (1978)., We also incorporate approximate temperature dependent limb darkening coefficients taken from \scite{Naimiy78}.495. Descriptions of previous versions of the code can be found in Shahbaz.Navlor&Charles(1993). and Somers.Mukai&Navlor (1996a)., Descriptions of previous versions of the code can be found in \scite{Shahbaz93} and \scite{Somers96a}.496. Phe model does not include the physical size ancl shape of the white dwarf or the white dwarl’s shadow onto the aceretion disc., The model does not include the physical size and shape of the white dwarf or the white dwarf's shadow onto the accretion disc.497 Also. it takes no account of a bright spot or an accretion stream.," Also, it takes no account of a bright spot or an accretion stream."498than the average spin teniperature othe eas means that the eas must have multiple phases. including a wari phase.,"than the average spin temperature of the gas means that the gas must have multiple phases, including a warm phase."499 Caven the relatively good agrcomment between the and 21-0 III coluun densities aud the reasonable value of the derived kinetic teniperature of the wide componcut. we consider it hielillikely that the side component is indeed gas iu the WNA phase.," Given the relatively good agreement between the and 21-cm HI column densities and the reasonable value of the derived kinetic temperature of the wide component, we consider it highly likely that the wide component is indeed gas in the WNM phase."500 Note that iu calculating the fraction of eas in the WNAL phase we have assuned that the Nyy imeasureineut is the best estimate of the total neutral-hbydrogen column clensity., Note that in calculating the fraction of gas in the WNM phase we have assumed that the $\NHI$ measurement is the best estimate of the total neutral-hydrogen column density.501 As discussed in the previous section. the 21-cin absorption profile of the +0.2212 DLA is in excellent aerecment with that expected were the absorption to arise in a inulti-phase 1ieciun similar to that of the Mls Way.," 	As discussed in the previous section, the 21-cm absorption profile of the $z=0.2212$ DLA is in excellent agreement with that expected were the absorption to arise in a multi-phase medium similar to that of the Milky Way."502 However. unlike the Calaxy. where the CNM and WNM both lads equitable contributions to the total HI colhuun density (Kulkarni&Ieiles 1988)). —75% ofthe neutral hivdroseu along this line ο [seht through the DLA umst be in the warn phase. ic. with temperature eSOOO Is. in order to account for its high estimated spin temperature.," However, unlike the Galaxy, where the CNM and WNM both make equitable contributions to the total HI column density \cite{kulkarni88}) ), $\sim 75\%$ of the neutral hydrogen along this line of sight through the DLA must be in the warm phase, i.e. with temperature $\sim 8000$ K, in order to account for its high estimated spin temperature."503 The average spin teniperature for the ;=0.2212 absorber of ~900 Ik is far higher than the typical spin teirperatures of 100—200 I& found in the Galaxy and nearby spirals (Braun&Walterbos1992.Draun 1997)).," 	The average spin temperature for the $z = 0.2212$ absorber of $\sim 900$ K is far higher than the typical spin temperatures of $100 - 200$ K found in the Galaxy and nearby spirals \cite{braun92, braun97}) )."504 Iheh spin temperatures of simular magnitude were earlier obtained by Coil et al. (, High spin temperatures of similar magnitude were earlier obtained by Carilli et al. (5051996) in DLAs at hieh redshift.,1996) in DLAs at high redshift.506 It was suggested there that the ligh Ty values at high : might be explained bv evolutionary effects., It was suggested there that the high $T_s$ values at high $z$ might be explained by evolutionary effects.507 Since then. however. there has been a substantial iucrease ia the wuuber of DLAs with 21-«4uu observations (Laneetal.1998.. Cheneahw&Kaucku1999.CheugalurIauckar2000.Ianekar&Chneusalur 2001)) and. as Cheugalur Ikanekar (2000) poiut out. high spin teniperatures appear to be typical for DLAs at all redslifts.," Since then, however, there has been a substantial increase in the number of DLAs with 21-cm observations \cite{lane98}, \cite{chengalur99, chengalur00,kanekar01}) ) and, as Chengalur Kanekar (2000) point out, high spin temperatures appear to be typical for DLAs at all redshifts."508 So ar. the ouly DLAs which show ow Ty values are those known to be associate with the disks of spiral galaxies.," So far, the only DLAs which show low $T_s$ values are those known to be associated with the disks of spiral galaxies."509 Cheugalur IKaueku (2000) (soe also Nanckar Choeugalur 2001) sugecsted that the higlh-T; DLAs were ikely to be associatc¢ with dwarf or LSD-tvpe galaxies. where a conibination of low central pressures. low dust conteut and low metallicities result in a simaller fraction for the ϱΝΑΙ," Chengalur Kanekar (2000) (see also Kanekar Chengalur 2001) suggested that the $T_s$ DLAs were likely to be associated with dwarf or LSB-type galaxies, where a combination of low central pressures, low dust content and low metallicities result in a smaller fraction for the CNM."510", Our conchision that ~A of the gas along the ine of siglt through the +=0.2212 DLA is iu the WNM yhase is dn good agreement with the observations of Young et al. (", Our conclusion that $\sim 75\%$ of the gas along the line of sight through the $z = 0.2212$ DLA is in the WNM phase is in good agreement with the observations of Young et al. (5112000). who find that nearby dwarf galaxies rave ονNUS of their neutral gas in the WNM phase.,"2000), who find that nearby dwarf galaxies have $\sim 80\%$ of their neutral gas in the WNM phase."512 Apart from possible selection effects. the association of DLAs with dwarf galaxies is somewhat surprising. both yecause of the existing paradiem of DLAs being associated with massive rotating disks (Prochaska&Wolfe1997.Prochaska&Wolfe 1998)). aud the expectation (based ou a ceusius of the III content of +=0 optically catalogued ealaxies by Rao&Briges1993)) that the bulk of the ueutral IIT at low redshifts is im large spiral galaxies.," 	Apart from possible selection effects, the association of DLAs with dwarf galaxies is somewhat surprising, both because of the existing paradigm of DLAs being associated with massive rotating disks \cite{prochaska97,prochaska98}) ), and the expectation (based on a census of the HI content of $z=0$ optically catalogued galaxies by \cite{rao93}) ) that the bulk of the neutral HI at low redshifts is in large spiral galaxies."513 Iloxwvever. blind searches for ΤΗ cussion at low redshift (ic. unbiased by the presence of a catalogued galaxy: Schneideretal.1998.Roseubere& 20003) also indicate that there could be substautial amounts of III in optically faint galaxies. although these results are still controversial (c.g. see Zwaan et al. (," However, blind searches for HI emission at low redshift (i.e. unbiased by the presence of a catalogued galaxy; \cite{schneider98, rosenberg00}) ) also indicate that there could be substantial amounts of HI in optically faint galaxies, although these results are still controversial (e.g. see Zwaan et al. ("5141997). for an opposing point of view).,"1997), for an opposing point of view)."515 Tn this context. it is of interest that the three lowest redshift DLAs known have all con. (tentatively) identified as dwrf (or LSB) galaxies (Bowenctal.2001.Turusheketal.2001.Cohen 20011) G@vhere. followine Turushek et al.," 	In this context, it is of interest that the three lowest redshift DLAs known have all been (tentatively) identified as dwarf (or LSB) galaxies \cite{bowen01,turnshek01,cohen01}) ) (where, following Turnshek et al."516" 2001. we only cousider svstems to be DLAs if they amect the ""classical selection criterion for a DLA. ie. for which an observed Lyius-oe profile yields an DIE column density Nyyο2«1077 23."," 2001, we only consider systems to be DLAs if they meet the “classical” selection criterion for a DLA, i.e. for which an observed $\alpha$ profile yields an HI column density $\NHI \ge 2\times 10^{20}$ )."517 Tu the case of the current absorber. Le Brun et al. (," In the case of the current absorber, Le Brun et al. ("5181997) sugeested that a galaxy at au impact parameter of G from OL 363 was likely to be the DLA lost.,1997) suggested that a galaxy at an impact parameter of $\sim 6^{''}$ from OI 363 was likely to be the DLA host.519 The spectrin of this ealaxy (Cohen2001)) confirms that it is dudeed at the correct redshift to produce the damped absorpion., The spectrum of this galaxy \cite{cohen01}) ) confirms that it is indeed at the correct redshift to produce the damped absorption.520 Turushek et al. (, Turnshek et al. (5212007) present both detailed uultiwaveleusth photometry aux optical spectroscopy of lis system. aud find that its colours are cousisCl with a cdwart (L~ VIL.) early-type galaxy.,"2001) present both detailed multi-wavelength photometry and optical spectroscopy of this system, and find that its colours are consistent with a dwarf $L \sim 0.1L_*$ ) early-type galaxy."522 The radial profile. jowever. indicates the presence of both a bulge aud a disk.," The radial profile, however, indicates the presence of both a bulge and a disk."523 Hence. these authors suggest that this ealaxy could be he equivalent of the dwiirf spirals seen at low redshifts (Schombertetal. 19051) aud/or might have evolved from he faint blue galaxies seen at ligher redshifts.," Hence, these authors suggest that this galaxy could be the equivalent of the dwarf spirals seen at low redshifts \cite{schombert95}) ) and/or might have evolved from the faint blue galaxies seen at higher redshifts."524 It should also be note that the iupact parameter (18) kpc is κοποματ large for a dwarf galaxy and it is thus also oossible that the absorption arises in an even fainter conipauiionu galaxy., It should also be noted that the impact parameter $\sim 18$ ) kpc is somewhat large for a dwarf galaxy and it is thus also possible that the absorption arises in an even fainter companion galaxy.525 The preseut detection of the WNAL in he; = absorber towards OI 363 is the second case of evidence for a multi-phase iuediuni in an extragalactic system., 	The present detection of the WNM in the $z=0.2212$ absorber towards OI 363 is the second case of evidence for a multi-phase medium in an extragalactic system.526 Lane et al. (, Lane et al. (5272000) found that the :=0.0912 absorber towards the same quasar also has a multiphase medium. with at most ~ one-third of the eas iu the ΝΑΙ phase.,"2000) found that the $z=0.0912$ absorber towards the same quasar also has a multi-phase medium, with at most $\sim$ one-third of the gas in the CNM phase."528 Interestingly. this DLA is also likely to be associated with a dwarf galaxy: Turushek ot al. (," Interestingly, this DLA is also likely to be associated with a dwarf galaxy; Turnshek et al. ("5292001) place an upper nuit of ~0.1L ouds I< baud huninosity.,2001) place an upper limit of $\sim 0.1~L_*$ on its K band luminosity.530 It is possible that such svclus dominate current samples of DLAs because the obscuration for lines of sight passing through gas with both ligh inetallicity and high dust couteut nuelt well be sufficient to make medi resolution spectroscopy ofthe backgrouu quasar extremely difficult (Fall&Pei1993 ), It is possible that such systems dominate current samples of DLAs because the obscuration for lines of sight passing through gas with both high metallicity and high dust content might well be sufficient to make medium resolution spectroscopy of the background quasar extremely difficult \cite{fall93}) ).531 Às conjectured earlier. the high observed spin teniperaures of DLAs thus seem to be a cousequence of their having a higher fraction of the WNM than is found for the Galaxy.," 	 As conjectured earlier, the high observed spin temperatures of DLAs thus seem to be a consequence of their having a higher fraction of the WNM than is found for the Galaxy."532 The fact that the two lowest redshift known DLAs have high ZF. uicaus that this cannot be due to evolutionary effects., The fact that the two lowest redshift known DLAs have high $T_s$ means that this cannot be due to evolutionary effects.533"to single out the effect of the stars, since that limit would remain even if the instrumental halo noise could be removed.","to single out the effect of the stars, since that limit would remain even if the instrumental halo noise could be removed."534" As an extreme example for a possible confusion error, it is worth mentioning that the orbit of the star S2 possibly was affected in 2002, during its pericenter passage, by such an event."," As an extreme example for a possible confusion error, it is worth mentioning that the orbit of the star S2 possibly was affected in 2002, during its pericenter passage, by such an event."535 Both recent analyses (??) therefore treat the respective 2002 data separately; either by ignoring it or by assigning large errors to it.," Both recent analyses \citep{Ghez:2008p945,Gillessen:2009p1117} therefore treat the respective 2002 data separately; either by ignoring it or by assigning large errors to it."536" In order to assess the magnitude of the confusion noise, we simulated stellar background populations in a Monte-Carlo fashion."," In order to assess the magnitude of the confusion noise, we simulated stellar background populations in a Monte-Carlo fashion."537 That needed two basic input distributions: a K-band luminosity function and the radial surface density profile., That needed two basic input distributions: a K-band luminosity function and the radial surface density profile.538 We based these on the findings of ?.., We based these on the findings of \cite{Genzel:2003p151}.539" We used three radial bins for our simulations: 0""«r10.2"", 0.2""c0.8"" and r3=3.5""."," We used three radial bins for our simulations: $0''<r_1<0.2''$, $0.2''<r_2<0.8''$ and $r_3=3.5''$."540" Since we had to assume the density also for stars much fainter than what can be measured, some assumptions had to be used."," Since we had to assume the density also for stars much fainter than what can be measured, some assumptions had to be used."541 For ra we extrapolated the cluster KLF of ? down to mx—24., For $r_3$ we extrapolated the cluster KLF of \cite{Genzel:2003p151} down to $m_\mathrm{K}=24$.542" For ri, and mx«18 we used the KLF as estimated from the S-stars cusp, scaled to the respective expected surface density from the radial profile."," For $r_{1,2}$ and $m_\mathrm{K} \le 18$ we used the KLF as estimated from the S-stars cusp, scaled to the respective expected surface density from the radial profile."543 At fainter magnitudes we extrapolated with a KLF from the same radial region that only counts stars which are not identified as late-type stars., At fainter magnitudes we extrapolated with a KLF from the same radial region that only counts stars which are not identified as late-type stars.544 This essentially assumes that for mx>18 only main sequence stars are present., This essentially assumes that for $m_\mathrm{K}>18$ only main sequence stars are present.545 Figure 13 shows the densities used., Figure \ref{f8} shows the densities used.546" Using the assumed densities per magnitude we simulated stellar fields, using Gaussian profiles with a FWHM of 30mas-42 mas(magnitudedependent, FWH Mofrealsourcesinthedeconvolvedy frames)Ogot =,cosV , and&binningof13 ο mas/pix."," Using the assumed densities per magnitude we simulated stellar fields, using Gaussian profiles with a FWHM of $30\,$ $-\,42\,$ mas (magnitude dependent, to mimic the FWHM of real sources in the deconvolved frames) and a binning of $13\,$ mas/pix."547"F oreachimage,ας. thetarg simulated images "," For each image, the target star was placed in the center of a box of 12 pixels width and background stars fainter than the target star were added."548per magnitude bin., The positional error of the target star is given by the difference between input position and the position at which it is found back by fitting the star with a Gaussian profile plus a floor.549 tomimicthe, We used up to $10^5$ simulated images per magnitude bin.550"The distribution of positional differences per magnitude bin was then fit with a Gaussian, the width of which estimates the position error."," The distribution of positional differences per magnitude bin was then fit with a Gaussian, the width of which estimates the position error."551 The resulting errors as a function of magnitude are shown in figure 13.., The resulting errors as a function of magnitude are shown in figure \ref{f8}.552" For the S-stars cluster (r< 0.8"") the error due to unrecognized confusion is of similar magnitude as the error due to halo noise.", For the S-stars cluster $r\lesssim0.8''$ ) the error due to unrecognized confusion is of similar magnitude as the error due to halo noise.553" For larger radii, the confusion induced error is smaller than the halo noise as the stellar densities drop rapidly with radius."," For larger radii, the confusion induced error is smaller than the halo noise as the stellar densities drop rapidly with radius."554" The increase of confusion error with stellar magnitude is well described by a power law of type Cx10°4""* (as the halo noise)."," The increase of confusion error with stellar magnitude is well described by a power law of type $C \times 10^{0.4\, m_\mathrm{K}}$ (as the halo noise)."555" Currently, relativistic effects have not yet been detected in the data of any star orbiting the GC MBH."," Currently, relativistic effects have not yet been detected in the data of any star orbiting the GC MBH."556" Actually, detecting 6? effects will be possible probably first in radial velocity measurements (?).."," Actually, detecting $\beta^2$ effects will be possible probably first in radial velocity measurements \citep{Zucker:2006p194}."557 Astrometrically detectable deviations from Newton's law have not yet been seen in the GC., Astrometrically detectable deviations from Newton's law have not yet been seen in the GC.558" Vice versa, the effects can currently be neglected in the analysis."," Vice versa, the effects can currently be neglected in the analysis."559" Still, a few effects are worth discussing here."," Still, a few effects are worth discussing here."560 We have analyzed a multitude of error sources that potentially influence and bias stellar positions as obtained from adaptive optics assisted imaging data in crowded, We have analyzed a multitude of error sources that potentially influence and bias stellar positions as obtained from adaptive optics assisted imaging data in crowded561"relie200 sinmlatious. respectively,","$relic200$ simulations, respectively."562 We choose to ouly fit halos above this παπά mass because at siialler scales additional plysics such as cooling not iucluded im our sinulatious would possibly stronely affect the emission., We choose to only fit halos above this minimum mass because at smaller scales additional physics such as cooling not included in our simulations would possibly strongly affect the emission.563 Additionally. we do not capture smedl mass halos that are Likely moving through these small clusters possibly creating a large fraction of the total radio enission.," Additionally, we do not capture small mass halos that are likely moving through these small clusters possibly creating a large fraction of the total radio emission."564 Because our simulation data does not have a measurable uncertainty for a given radio power. we have to use an alternate method of determining the error estinates of our paranueters.," Because our simulation data does not have a measurable uncertainty for a given radio power, we have to use an alternate method of determining the error estimates of our parameters."565 We first fud the best fit parameters using a uniform weighting., We first find the best fit parameters using a uniform weighting.566 By caleulatiug the residuals for cach point from this best-fit relation. we estimate the muiform error for cach point as the standard deviation of this residual.," By calculating the residuals for each point from this best-fit relation, we estimate the uniform error for each point as the standard deviation of this residual."567 We then ft the data again using this error to obtain the uncertainty estimates in cach parameter., We then fit the data again using this error to obtain the uncertainty estimates in each parameter.568 The values of these parameters are shown inTable 1.., The values of these parameters are shown inTable \ref{tab:fitpars}. .569Amendola L. D. Tocehini-Valentini. astro-ph/0111535.Phys.,"Amendola L. D. Tocchini-Valentini, astro-ph/0111535, Phys."570" Rev.D66.. 043!cVTble Amendola L.. C. Quercellini. D. Tocehini-Valentini and. A. ακοή, (2002) ph/0205097 Amendola L.. M. Gasperini. D. Tocchini-Valentini and C. Unearelli. (2002) Baccigalupi C...A. Balbi. ο, Matarrese. F. Perrotta. N. Vittorio. astro-ph/0109007. Phys."," Rev., 043528 Amendola L., C. Quercellini, D. Tocchini-Valentini and A. Pasqui, (2002) astro-ph/0205097 Amendola L., M. Gasperini, D. Tocchini-Valentini and C. Ungarelli, (2002) Baccigalupi C., A. Balbi, S. Matarrese, F. Perrotta, N. Vittorio, astro-ph/0109097, Phys."571Rev. (2002) Dahcall N. et al.,Rev. (2002) Bahcall N. et al.572 2002. Dean R. Melehiorri A.. astro-ph/0110472. Phys.," 2002, Bean R. Melchiorri A., astro-ph/0110472, Phys."573Rev. (2002) Benilez N. et al.. (,"Rev. (2002) Benitez N. et al., ("5742002) Bonanno A. M. Reuter (2002) Phys.,2002) Bonanno A. M. Reuter (2002) Phys.575 Lett., Lett.576 D 527. Caldwell BR... Dave R. Steinhardt DJ. (1993). Phys.," B 527, Caldwell R.R., Dave R. Steinhardt P.J. (1998), Phys."577 Rev. Lett., Rev. Lett.578 80. 1582 Carvalho J.C... J.A.S. Lima. Ll. Waga. Rev.D46::2404-2407.1992 Chimento L.P.. A. $. Jakubi D. Pavon. Phys.," 80, 1582 Carvalho J.C., J.A.S. Lima, I. Waga, :2404-2407,1992 Chimento L.P., A. S. Jakubi D. Pavon, Phys."579 Rev.D62.. OS (2000). Corasaniti E. Copeland. Phys.," Rev., 063508 (2000), astro-ph/0005070; Corasaniti E. Copeland, Phys."580Hev. (2002) Damour Esposito-Farese G.. Class.,"Rev. (2002) Damour Esposito-Farese G., Class."581 Quantum Gray., Quantum Grav.582 9. 2093 Damour T. Nordtvedt Ix.. Plivs.," 9, 2093 Damour T. Nordtvedt K., Phys."583 Rev. Lett., Rev. Lett.584 70. 2217 Damour T.. G. W. Gibbons and C. Gundlach. Phys.," 70, 2217 Damour T., G. W. Gibbons and C. Gundlach, Phys."585 Rev. Lett..," Rev. Lett.,"586 64. 123. Dalal N. et al..," 64, 123, Dalal N. et al.,"587 Phys., Phys.588 Rev. Lett..," Rev. Lett.,"58987.. 141802 De Bernarclis et al., 141302 De Bernardis et al.590Nature, Nature591lhuninostv iu the range 1l«1007l0Peres +.,"luminosity in the range $\rm 1\times 10^{42} < Ly\alpha < 5921.5 \times 10^{43} erg \; s^{-1}$ ."593 Here we assumed that the LLF docs not evolve from :—6.6 to :—7.1., Here we assumed that the LF does not evolve from $z$ =6.6 to $z$ =7.7.594 Each galaxy was then assigned a raudonm redshift τε«tolg where ze andl zy correspoud to the minima and imaxiuuui wavelengths where the trausuüssion of the UND filter drops to zero., Each galaxy was then assigned a random redshift $z_{L} < z < z_{H}$ where $z_{L}$ and $z_{H}$ correspond to the minimum and maximum wavelengths where the transmission of the UNB filter drops to zero.595" Next. to each galaxy we assigned a flux F=Lis,Ind; where dp is the luminosity distance."," Next, to each galaxy we assigned a flux $\rm F=L_{Ly\alpha}/4\pi d_{L}^{2}$ where $\rm d_{L}$ is the luminosity distance."596 We distribute this flux in waveleneth using an asviunietrie hue profile drawn from the +=5.7 spectra of Rhoadsetal. (2003).., We distribute this flux in wavelength using an asymmetric line profile drawn from the $z=5.7$ spectra of \citet{rho03}. .597 The flux transiuitted through the UND filter was then determined as feesffxTadd (where Ty is the filter transmission and fy the fux density of the cluission line).," The flux transmitted through the UNB filter was then determined as $f_{trans} = \int f_\lambda T_\lambda 598d\lambda$ (where $T_\lambda$ is the filter transmission and $f_\lambda$ the flux density of the emission line)."599 This accounts for the loss of the flux that results from a filter whose width is comparable to the line width (and not much greater as would be the case for a filter)., This accounts for the loss of the flux that results from a filter whose width is comparable to the line width (and not much greater as would be the case for a filter).600 We then created a histogram of magnitudes after converting the couvolved fiux to magnitudes calculated using the following relation: and with e the speed of light., We then created a histogram of magnitudes after converting the convolved flux to magnitudes calculated using the following relation: and with $c$ the speed of light.601 Lastly. to include the instrumental effects. we multiplied the ΡΟ of galaxies in cach maeuitude biu bv the corresponding recovery fraction obtained frou our artificial source simulations in our UND image(sece section 2.1).," Lastly, to include the instrumental effects, we multiplied the number of galaxies in each magnitude bin by the corresponding recovery fraction obtained from our artificial source simulations in our UNB image(see section 2.4)."602 We then converted cach magnitude bin to a ]huunünositv biu. and counted the nuuber of detected ealaxies iu cach hunuinositv biu.," We then converted each magnitude bin to a luminosity bin, and counted the number of detected galaxies in each luminosity bin."603 We repeated this simulation teu times. aud taking au average. we found that about one eenitter should be expected im our survey.," We repeated this simulation ten times, and taking an average, we found that about one emitter should be expected in our survey."604 It should be noted that we assumed a nou-evolving LLF from :26.6 to τιfadτν and that every ecnmütter has the same asviuuetrie line profile.," It should be noted that we assumed a non-evolving LF from $z$ =6.6 to $z$ =7.7, and that every emitter has the same asymmetric line profile."605 While we expect about one eenuütter iun our survey there are huge uncertainties mainly due to the Poisson noise. aud field to field variation or cosmic variance.," While we expect about one emitter in our survey there are large uncertainties mainly due to the Poisson noise, and field to field variation or cosmic variance."606 Tibietal.(20090) have estimated feld to field variation of οσους to be 230% for a volume and flux limited ssurvey with a survey volue ~2100Mpc?., \citet{til09} have estimated field to field variation of emitters to be $\gtrsim 30\%$ for a volume and flux limited survey with a survey volume $\rm \sim 2 \times 10^{5} \; Mpc^{3}$.607 We expect a larger field to field variation for suialler survey volutes., We expect a larger field to field variation for smaller survey volumes.608 We also estimated the cosmic variance expected im our survey using the cosmic variance caleulator (Trenti&Sti-avelli 20083., We also estimated the cosmic variance expected in our survey using the cosmic variance calculator \citep{tre08}.609. For our survey we should expect à cosnic variance of about 5854 assuming au intrinsic uuuber of ssourees at lo=fad in agreement with a non-evolviug LLF from :=6.6 (Iwashikawaetal.2006) to --7.7., For our survey we should expect a cosmic variance of about $58\%$ assuming an intrinsic number of sources at $z=7.7$ in agreement with a non-evolving LF from $z=6.6$ \citep{kas06} to $z=7.7$.610 On the other hand our caudidate counts are quite consistent with the buninosity function at z=5.7 2004)).., On the other hand our candidate counts are quite consistent with the luminosity function at z=5.7 \citep{ouc09}. .611 Using a large sample of candidates. (2008) fouud no aut evolution of LLF between : siguific Land :—5.7.," Using a large sample of candidates, \citet{ouc08} found no significant evolution of LF between $z$ =3.1 and $ z$ =5.7."612 The evolution of the LLF between 2 —5.7 aud :=6.5 is not conclusive., The evolution of the LF between $z=$ 5.7 and $z=6.5$ is not conclusive.613 For exaniple. Malliotra&Rhoads(2001) found no significant evolution of LLFbetween 2=5.7 aud :=6.5. while IashikawactinH.(2006) sugeest an evolution of bright cud of the LE this redshift range.," For example, \citet{mr04} found no significant evolution of LF between $z$ =5.7 and $z=6.5$, while \citet{kas06} suggest an evolution of bright end of the LF in this redshift range."614 Ou the theoretical frout. several models(Thomunes&Aeiscnheiner2005:Furlanettoot9) havebeen developed to predict redshift evolutiou of the LLF.," On the theoretical front, several models \citep{tho05,fer05, del06, dij07,kob07,mcq07, day08, nag08,sam09,til09} have been developed to predict redshift evolution of the LF."615 While several imiodoels (6.8.Sammictal.2009:etal.2009) predict no significant evolution of LLF at :€7. the predictions differ greatly amoug different models.," While several models \citep[e.g.][]{sam09, til09} predict no significant evolution of LF at $z\lesssim 7$, the predictions differ greatly among different models."616 These differeuces among the models can be attributed to differiug input assumptions. which in turn stem from our iniperfect undoerstaudius of the physical nature of aoeealaxies. aud from the simall samples currently availableat lieh redshift.," These differences among the models can be attributed to differing input assumptions, which in turn stem from our imperfect understanding of the physical nature of galaxies, and from the small samples currently availableat high redshift."617 At 2> 6.5. there are ouly a few searches for celitters.," At $z>$ 6.5, there are only a few searches for emitters."618 Iveetal.(2006). found. one spectroscopically confirmed LAE at :=6.96. audcurrently there are no spectroscopically confirmed LAEs at :> ," \citet{iye06} found one spectroscopically confirmed LAE at $z$ =6.96, andcurrently there are no spectroscopically confirmed LAEs at $z>$ 7."619However. there are few photometric searches for," However, there are few photometric searches \citep{par94, wil05, cub07, hib09} for"620Observations of Active Galactic Nuclei (AGN) have revealed tha many of them are partially obscured by material in our line of sight. within the inner tens of parsecs of the central engine (e.g. ? and ? or a review).,"Observations of Active Galactic Nuclei (AGN) have revealed that many of them are partially obscured by material in our line of sight, within the inner tens of parsecs of the central engine (e.g. \citealt*{risaliti99} and \citealt*{maiolino&risaliti07} for a review)."621" This obscuring material will influence both the AG and our observations. and the study of its properties is fundamenta or an unbiased understanding o"" AGN physics."," This obscuring material will influence both the AGN and our observations, and the study of its properties is fundamental for an unbiased understanding of AGN physics."622 The gas arounc he nucleus will be under the efect of the inward gravitationa orce of the supermassive black 10le and the outward pressure of he radiation emitted in the cenral region., The gas around the nucleus will be under the effect of the inward gravitational force of the supermassive black hole and the outward pressure of the radiation emitted in the central region.623 By investigating the balance between these two forces. one can predict the behaviour of the gas.," By investigating the balance between these two forces, one can predict the behaviour of the gas."624" The Eddington luminosity. £p. is defined as the value at which the radiation pressure balances the gravitational force of he black hole: £=Ly. with Le=4taCimedlyfor. where G is the gravitational constant. m, the proton mass. ο the speed of ight. Mig the mass of the black hole and op the cross-section for Thomson scattering."," The Eddington luminosity, $L_{\rm E}$, is defined as the value at which the radiation pressure balances the gravitational force of the black hole: $L = L_{\rm E}$, with $L_{\rm E}=4\pi Gm_{\rm p}cM_{\rm BH}/\sigma_{\rm T}$, where G is the gravitational constant, $_{\rm p}$ the proton mass, c the speed of light, $M_{\rm BH}$ the mass of the black hole and $\sigma_{\rm T}$ the cross-section for Thomson scattering."625 The Eddington ratio (A=£/ Lp). is then a measure of the balance between these two forces. for a certain AMyg.," The Eddington ratio $\lambda = L/L_{\rm E}$ ), is then a measure of the balance between these two forces, for a certain $M_{\rm BH}$."626 In the presence of dust. the gas couples with the dust grains via Coulomb interactions and the cross-section for the interaction with photons is considerable enhanced.," In the presence of dust, the gas couples with the dust grains via Coulomb interactions and the cross-section for the interaction with photons is considerable enhanced."627 The effective cross-section for dusty gas can be defined as 4=shor with a boost factor .1 (2.. 2.. 2).," The effective cross-section for dusty gas can be defined as $\sigma_{\rm d}=A\sigma_{\rm T}$ with a boost factor $A$ \citealt{fabiancelottierlund06}, \citealt*{fabian&vasudevan08}, \citealt{fabian09}) )."628 We can then determine an effective Eddington ratio for dusty gas as a function of the classical Eddington ratio: It follows that. A= L/-Lis now the limit at which the radiation pressure from the black hole is able to expel the mass of dusty gas around it.," We can then determine an effective Eddington ratio for dusty gas as a function of the classical Eddington ratio: It follows that, $\lambda = 1/A$ is now the limit at which the radiation pressure from the black hole is able to expel the mass of dusty gas around it."629 2 and ? explored the effective Eddington limit for dusty gas by investigating the properties of AGN samples in the Deep Field South (CDF-S). Lockman Hole and in the local Universe. and found that the objects tend to lie below their effective Eddington limit. as expected.," \cite{fabian&vasudevan08} and \cite{fabian09} explored the effective Eddington limit for dusty gas by investigating the properties of AGN samples in the Deep Field South (CDF-S), Lockman Hole and in the local Universe, and found that the objects tend to lie below their effective Eddington limit, as expected."630 AGN are expected to interact with the surrounding galaxy. affecting its evolution.," AGN are expected to interact with the surrounding galaxy, affecting its evolution."631 In fact. there is evidence of a close," In fact, there is evidence of a close"632as a random realization of some DF. then each value calculated in this system can be considered as a random variable.,"as a random realization of some DF, then each value calculated in this system can be considered as a random variable."633 Let Q; be random variable defined as the value Q calculated in piece i., Let $Q_i$ be random variable defined as the value $Q$ calculated in piece $i$.634 For example. if the two systems under consideration are just two random realization of the same DF. then 4; and q»; are two samples of the random variable Q;.," For example, if the two systems under consideration are just two random realization of the same DF, then $q_{1,i}$ and $q_{2,i}$ are two samples of the random variable $Q_i$."635 We can estimate the variance of this random variable., We can estimate the variance of this random variable.636 Let var); and var»; be estimates of the variance of Q; calculated for the first and the second system. respectively. and let us consider the value (KenneyandKeeping.1951).. (KenneyandKeeping.1951.p.164).. (AT))," Let $var_{1,i}$ and $var_{2,i}$ be estimates of the variance of $Q_i$ calculated for the first and the second system, respectively, and let us consider the value \citep{K51}. \citep[p. 164]{K51}. \ref{eq_c})"637"Cy; (solid). c,,; (dotted). aud s (dashed).","$e_{y,z}$ (solid), $e_{y,z}$ (dotted), and $s$ (dashed)."638 Note that Chu(Ng). and ον(kg). are zero. for ky because of the periodic boundaries.," Note that $e_{z,y}(k_0)$ and $e_{z,x}(k_0)$ are zero for $k_0$ because of the periodic boundaries."639" The shear term s primarily drives the variation of |B,(hy)P. fippiug sign as [B.Ahy)? goes to zero."," The shear term $s$ primarily drives the variation of $|\tilde{B}_y(k_0)|^2$, flipping sign as $|\tilde{B}_x(k_0)|^2$ goes to zero."640 In contrast. ον. 1s generally negative. acting as turbulent resistivity.," In contrast, $e_{x,z}$ is generally negative, acting as turbulent resistivity."641" The ος(0) term is more crratic, frequently flipping sign over a sinele cycle. but the net effect is an overall oscillation of [D(&9)|? over ~7 orbital periods."," The $e_{y,z}(k_0)$ term is more erratic, frequently flipping sign over a single cycle, but the net effect is an overall oscillation of $|\tilde{B}_x(k_0)|^2$ over $\sim 7$ orbital periods."642 Tu many respects. the behavior we see in the stratified «λατοις is simular to that observed in the uustratified. zoro-uet flux calculations of Lesur&Ogilvie(2008).," In many respects, the behavior we see in the stratified simulations is similar to that observed in the unstratified, zero-net flux calculations of \citet{lo08}."643 Using au incompressible spectral code. they find dynamo evcles with a ~5 orbit periodicity.," Using an incompressible spectral code, they find dynamo cycles with a $\sim 5$ orbit periodicity."644 This is similar to the oscillatious in our stratified runs where rms power on large scales is broadly distributed ou times scales ~610 orbits (see Figure 7))., This is similar to the oscillations in our stratified runs where rms power on large scales is broadly distributed on times scales $\sim 6-10$ orbits (see Figure \ref{f:kvsf}) ).645 The normalized quantities plotted in Figure 17. are to them Equation (16)., The normalized quantities plotted in Figure \ref{f:emft} are to their Equation (16).646 Comparison of Figure 17 with Figs, Comparison of Figure \ref{f:emft} with Figs.647 [| aud 5 in their paper. show that the behavior of the EMES during oscillations are also quite simular. sugeesting that a common (or. at least. related) niechanisni iav be respousible for these oscillations.," 4 and 5 in their paper, show that the behavior of the EMFs during oscillations are also quite similar, suggesting that a common (or, at least, related) mechanism may be responsible for these oscillations."648 This motivates a more detailed comparison between uustratified and stratified ruus in future work., This motivates a more detailed comparison between unstratified and stratified runs in future work.649 Although it is useful to focus ou the vertical wave vectors when trving to understand properties of large scale fields. an understanding of the overall powcr spectrum beuefits from an analysis of the shell integrated quantities.," Although it is useful to focus on the vertical wave vectors when trying to understand properties of large scale fields, an understanding of the overall power spectrum benefits from an analysis of the shell integrated quantities."650 We plot the time aud shell average EMES terms in Figure 18. for S32R1Z1 (black) S6LRIZL (blue). and SI28RIZLE (red).," We plot the time and shell average EMFs terms in Figure \ref{f:emf} for S32R1Z4 (black) S64R1Z4 (blue), and S128R1Z4 (red)."651 As iu Figure 17.. these quantities are normalized by the shell inteerated power spectra.," As in Figure \ref{f:emft}, these quantities are normalized by the shell integrated power spectra."652 Each normalized term is then time averaged frou 50-300 orbits., Each normalized term is then time averaged from 50-300 orbits.653 Since the amplitudes of the maguctic energv densities ποσα to be in statistical steady states over this period. we prestune the left haud sides of (12)) and (13)) are nearly zero.," Since the amplitudes of the magnetic energy densities seem to be in statistical steady states over this period, we presume the left hand sides of \ref{eq:fourx}) ) and \ref{eq:foury}) ) are nearly zero."654 Therefore the sum of the terms in each panel must be balanced by uuucerical dissipation ternis. as discussed iu previous work (Fromang&Papaloizou," Therefore the sum of the terms in each panel must be balanced by numerical dissipation terms, as discussed in previous work \citep{fp07,shb09}."655" Iu the top panel we plot the terms οςμή) (solid) and ορ) (dotted) which coutribute to evolution of D,."," In the top panel we plot the terms $e_{z,y}(k)$ (solid) and $e_{y,z}(k)$ (dotted) which contribute to evolution of $B_x$."656 At the large scales. we find that the c. terii is more inrportant for Ποια ecucration and its normalized amplitude is nearly independent of resolution.," At the large scales, we find that the $e_{z,y}$ term is more important for field generation and its normalized amplitude is nearly independent of resolution."657" The «,,.- termi is simaller iu amplitude and slightly negative as large scales."," The $e_{y,z}$ term is smaller in amplitude and slightly negative as large scales."658 Even though ορ. teuds to oscillates about zero over an individual dynamo cvele while ος ustially renmuus positive. the amplitude of ος 38 generally larecr. so the dominance of e; at large scales isnot sinply the result of time averaging.," Even though $e_{y,z}$ tends to oscillates about zero over an individual dynamo cycle while $e_{z,y}$ usually remains positive, the amplitude of $e_{z,y}$ is generally larger, so the dominance of $e_{z,y}$ at large scales is simply the result of time averaging."659" As one moves to smaller scales. ο Visesand eventually dominates the geucration of B,.."," As one moves to smaller scales, $e_{y,z}$ risesand eventually dominates the generation of $B_x$."660 The characteristic vaveuumber at which the crossing occurs shifts to higher values as the resolution increases., The characteristic wavenumber at which the crossing occurs shifts to higher values as the resolution increases.661" The bottom. panel shows ¢..(4) (solid). ον.) (dotted) aud s (dashed). the terms which contribute growth in B,."," The bottom panel shows $e_{z,x}(k)$ (solid), $e_{x,z}(k)$ (dotted) and $s$ (dashed), the terms which contribute growth in $B_y$."662" At large scales growth of B, is dominated bv the shear term. while both 6. aud ορ. are of comparable magnitude and negative."," At large scales growth of $B_y$ is dominated by the shear term, while both $e_{z,x}$ and $e_{x,z}$ are of comparable magnitude and negative."663" At small scales. €, and ος both grow. becoming positive and dominating over the shear term."," At small scales, $e_{z,x}$ and $e_{x,z}$ both grow, becoming positive and dominating over the shear term."664 Again. the wavemmuber of the crossover Hiereases with resolutio1.," Again, the wavenumber of the crossover increases with resolution."665 Authors have often focused oi horizontally average properties of the flow or (equivaleutlv) the power spectral variation onlv along vertical wav| vectors (οιFromanediscussed above)..," Authors have often focused on horizontally average properties of the flow or (equivalently) the power spectral variation only along vertical wave vectors \citep[e.g][ which were discussed666above]{fp07,lo08}."667 We note that the behavior of ομ.. aud exy we lave described diffCrs slenificantly frou what one would iufer if ouly verical wavevectors were considered.," We note that the behavior of $e_{y,z}$ and $e_{z,y}$ we have described differs significantly from what one would infer if only vertical wavevectors were considered."668 As previously mcutiored. the svuuuetries of the periodic box force ονμή.) to be zero. and oulv Cyzv) contributes.," As previously mentioned, the symmetries of the periodic box force $e_{z,y}(k_z)$ to be zero, and only $e_{y,z}(k_z)$ contributes."669 However. it is clear from Figure Ls that the vertical EMF and its toroidal variation is also esseutial for uuderstaudiug the mechanism which sustains turbulence in these simulations.," However, it is clear from Figure \ref{f:emf} that the vertical EMF and its toroidal variation is also essential for understanding the mechanism which sustains turbulence in these simulations."670 The question remains as to why the addition. of stratification leads to couvergence iu the turbulent stresses and cherey densities., The question remains as to why the addition of stratification leads to convergence in the turbulent stresses and energy densities.671 One possibility is that development of local toroidal field is key to sustainiug turbulence iu both stratified and wustratified domains., One possibility is that development of local toroidal field is key to sustaining turbulence in both stratified and unstratified domains.672 It is possible that the strength of toroidal field is eutirely set bv the resolution in unstratified domains. while stratified domains offer a characteristic scale which is independent of resolution. duc to the action of the large scale dvnamo.," It is possible that the strength of toroidal field is entirely set by the resolution in unstratified domains, while stratified domains offer a characteristic scale which is independent of resolution, due to the action of the large scale dynamo."673 Indeed. it has already been demonstrated (e.g.Hawleyetal.1995:Sinon&2009) that a elobal net toroidal field leads to enhanced turbulent enerev densitics aud stresses. aud leads to couvergeuce in the stress as resolution inereses (Caranetal. 2009)..," Indeed, it has already been demonstrated \citep[e.g.][]{hgb95,sh09} that a global net toroidal field leads to enhanced turbulent energy densities and stresses, and leads to convergence in the stress as resolution increses \citep{gua09}. ."674 Furthermore. our sinmlatious show a correlation between the stress and the streugth of the mean toroidal field. both elobally in the two scale height averages (Table," Furthermore, our simulations show a correlation between the stress and the strength of the mean toroidal field, both globally in the two scale height averages (Table"675A trivial solution of equation (5)) is both .X=laud Y—Lin the meridional plane wherein the perturbing object lies (42=0). which indicates the (vertically) highest intersecting points between the outer and inmer boundaries. le.. the extension limit of the intensively perturbed region.,"A trivial solution of equation \ref{equ:mor}) ) is both $X=1$ and $Y=1$ in the meridional plane wherein the perturbing object lies $\varphi=0$ ), which indicates the (vertically) highest intersecting points between the outer and inner boundaries, i.e., the extension limit of the intensively perturbed region."676 X=1 aud Y—1 staud for au ellipse aud a hyperbola. by definition. with the common foci at the location of the perturber and the mirror point about the orbital axis.," $X=1$ and $Y=1$ stand for an ellipse and a hyperbola, by definition, with the common foci at the location of the perturber and the mirror point about the orbital axis."677" However. both .X=1 aud Y=1 actually correspoud to the same formula for a hyperbola since we only consider the shape of a wake induced by a supersonic perturber (M4,>1)."," However, both $X=1$ and $Y=1$ actually correspond to the same formula for a hyperbola since we only consider the shape of a wake induced by a supersonic perturber $\mach>1$ )."678" The vertex of the hyperbola. r/ry=Mj,! correspouds to the distance of the junction between the outer aud inner arm boundaries in the orbital plane."," The vertex of the hyperbola, $r/r_p=\mach^{-1}$, corresponds to the distance of the junction between the outer and inner arm boundaries in the orbital plane."679" The extension limit of the high density ares in the vertical direction is simplified to the asyimptotes of the lyperbola. which are overlaid by black lines in Figure Lbb. This tuclicates that the angular size of the ares is 21anον—13E?, further extended with higher orbital Mach number."," The extension limit of the high density arcs in the vertical direction is simplified to the asymptotes of the hyperbola, which are overlaid by black lines in Figure \ref{fig:ptm}b b. This indicates that the angular size of the arcs is $2\tan^{-1}(\mach^2-1)^{1/2}$, further extended with higher orbital Mach number."680" This vertical stretch. however. does not depend very strougly on «τρ except for near unity (1.e.. 607-— (lor M,z 2)."," This vertical stretch, however, does not depend very strongly on $\mach$ except for near unity (i.e., – for $\mach\ge2$ )."681 Iu Figure Lbb iudividual segments of the boundaries. satisfying equation (5)). appear to be circular arcs about the centers at either the object position or the mirror point.," In Figure \ref{fig:ptm}b b individual segments of the boundaries, satisfying equation \ref{equ:mor}) ), appear to be circular arcs about the centers at either the object position or the mirror point."682" In the plane of the object (422 Orr 0).the ares extended from the outer boundaries of the spiral arm in Figure laa (Le. r/ry=1. 1-222MI. b+ME,T 14+Gs MET. 1+80M, T. and 1+108PpE at z/ry 0) center around the perturbing object at. (e.2)/rg=(1.0)."," In the plane of the object $\varphi=0$ ; $x>0$ ),the arcs extended from the outer boundaries of the spiral arm in Figure \ref{fig:ptm}a a (i.e., $x/r_p=1$, $1+2\pi\mach^{-1}$, $1+4\pi\mach^{-1}$, $1+6\pi\mach^{-1}$ $1+8\pi\mach^{-1}$ , and $1+10\pi\mach^{-1}$ at $z/r_p=0$ ) center around the perturbing object at $(x,z)/r_p=(1,0)$."683 These concentric arcs have endpoints on the hyperbola defined in equation (9)) (or. the asymptotes in eq. [10]]).," These concentric arcs have endpoints on the hyperbola defined in equation \ref{equ:hyp}) ) (or, the asymptotes in eq. \ref{equ:asy}] ]),"684" except for the point (not visible) at the perturber position and the curve passing througl ο...25M,LO). which completes its full circle."," except for the point (not visible) at the perturber position and the curve passing through $(x,z)/r_p=(1+2\pi\mach^{-1},0)$, which completes its full circle."685" The inner :-intercept of this circle at /ry=1—23M,! actually corresponds to the inner edge of the curved Mach cone. which is cousidered as a part of the outer arm boundary (see relsec:spi)) aud. here. tliis is clarified."," The inner $x$ -intercept of this circle at $x/r_p=1-2\pi\mach^{-1}$ actually corresponds to the inner edge of the curved Mach cone, which is considered as a part of the outer arm boundary (see \\ref{sec:spi}) ) and, here, this is clarified."686 Ou the other haud. the arcs crossing the inner arin boundaries on this side Gr> 0) have the center at the mirror point (e.2)/ry=(71.0). resulting in larger radii compared to the radii of the outer counterparts.," On the other hand, the arcs crossing the inner arm boundaries on this side $x>0$ ) have the center at the mirror point $(x,z)/r_p=(-1,0)$, resulting in larger radii compared to the radii of the outer counterparts."687" The outer boundaries ou the opposite side Gr« 0). Le.. the ares having intercepts at ur/ry=—1—aM,| -]1—SUME,| -1-onME, οDOM,B aud —1-—Us.ME,. are concentric about the mirror point of the object. Gr.2)/ry=(-1.0)."," The outer boundaries on the opposite side $x<0$ ), i.e., the arcs having intercepts at $x/r_p=-1-\pi\mach^{-1}$, $-1-3\pi\mach^{-1}$, $-1-5\pi\mach^{-1}$, $-1-7\pi\mach^{-1}$, and $-1-9\pi\mach^{-1}$, are concentric about the mirror point of the object, $(x,z)/r_p=(-1,0)$."688 On this side Ge« 0). iustead. the center of the inner boundaries is at the objectposition. Cr.z)/ry= (1.0).," On this side $x<0$ ), instead, the center of the inner boundaries is at the objectposition, $(x,z)/r_p=(1,0)$ ."689 The inner aud outer boundaries. in fact. constitute the same circles. which is better visualized in Figure 2. for the number of perturbations.," The inner and outer boundaries, in fact, constitute the same circles, which is better visualized in Figure \ref{fig:num} for the number of perturbations."690 Particularly.the largest sphere about," Particularly,the largest sphere about"691"very long ""filaments? in the bottom panel of Figure 10..",very long “filaments” in the bottom panel of Figure \ref{fig:S4densityglobus}.692 This cluster is a remnant οἱ cloud 1. and its mass is ~3000 solar Dlasses.," This cluster is a remnant of cloud 1, and its mass is $\sim 3000$ solar masses."693 In fact most of the stars in simulations $83 and 84 belong to a cluster., In fact most of the stars in simulations S3 and S4 belong to a cluster.694 This clustered mode of star formation in the fast-cooling runs is due to the gas collapsing promptly and thus forming very massive dense gas halos., This clustered mode of star formation in the fast-cooling runs is due to the gas collapsing promptly and thus forming very massive dense gas halos.695 We suspect that these results would have been dillerent had. we been able to treat radiative feedback from star formation in our simulations., We suspect that these results would have been different had we been able to treat radiative feedback from star formation in our simulations.696 We note that in such a dense environment. any feedback from star formation activity. such as radiation or outflows. would not have been able to escape easily and would therefore. have heated the surrounding. gas. suppressing further fragmentation and increasing the Jeans mass (??7).," We note that in such a dense environment, any feedback from star formation activity, such as radiation or outflows, would not have been able to escape easily and would therefore have heated the surrounding gas, suppressing further fragmentation and increasing the Jeans mass \citep{Nayakshin06a,KrumholzEtal07}."697. Rather than the multiple low-mass stars in a cluster we would perhaps have seen fewer but higher mass stars., Rather than the multiple low-mass stars in a cluster we would perhaps have seen fewer but higher mass stars.698 The mass spectrum of stars would. probably also change if the colliding clouds had net rotation. (spin) or a turbulent structure before the collision. providing some stability against collapse.," The mass spectrum of stars would probably also change if the colliding clouds had net rotation (spin) or a turbulent structure before the collision, providing some stability against collapse."699 All these effects are however minor as far as orbital motion of the stars is concerned. which is the focus of our paper.," All these effects are however minor as far as orbital motion of the stars is concerned, which is the focus of our paper."700 Simulation $4 is the fast-cooling equivalent of 82., Simulation S4 is the fast-cooling equivalent of S2.701 As such. the inner gaseous disc in S4 is subject to a similar intermittent infall of gas [rom larger radii.," As such, the inner gaseous disc in S4 is subject to a similar intermittent infall of gas from larger radii."702 The newly arriving gas tweaks the disc orientation signilicantly. causing it to undergo midplane rotation (refer back to ?7)).," The newly arriving gas tweaks the disc orientation significantly, causing it to undergo midplane rotation (refer back to \ref{sec:s2dynamics}) )."703 The cdillerence here however is that disc is fragmenting into stars whilst its orientation is changing., The difference here however is that disc is fragmenting into stars whilst its orientation is changing.704 Interestingly. these stars do not follow the evolving disc orientation but rather remain in their original configuration.," Interestingly, these stars do not follow the evolving disc orientation but rather remain in their original configuration."705" This is clearly seen in Figure 11.. where the stellar. disc remembers the ""old"" orientation of the gaseous disc in which the stars were born. whereas the gaseous. disc evolves to quite a different orientation."," This is clearly seen in Figure \ref{fig:S4_inner_disk}, where the stellar disc remembers the “old” orientation of the gaseous disc in which the stars were born, whereas the gaseous disc evolves to quite a different orientation."706 The possibility of this effect taking place was suggested by ?).., The possibility of this effect taking place was suggested by \cite{NC05}.707 Phese authors found that stars will not follow clisc midplane changes if these occur faster than the “critical rotation time”. which is estimated to be about 500 code units lor the disc in Figure 11..," These authors found that stars will not follow disc midplane changes if these occur faster than the “critical rotation time”, which is estimated to be about 500 code units for the disc in Figure \ref{fig:S4_inner_disk}."708 Our simulations are thus consistent with these predictions as the disc orientation changed on a timescale of just SO code units., Our simulations are thus consistent with these predictions as the disc orientation changed on a timescale of just $80$ code units.709 As a result of a several star forming events in gaseous discs of dilferent orientations. the inner stellar disc in 84 is very much different from that found in S2. and even more so when compared to S1.," As a result of a several star forming events in gaseous discs of different orientations, the inner stellar disc in S4 is very much different from that found in S2, and even more so when compared to S1."710 In the latter cases. when cooling is more gradual (io. ο)= 1). stars only form in the inner gaseous disc once it has settled into a relatively stable orientation.," In the latter cases, when cooling is more gradual (i.e. $\beta=1$ ), stars only form in the inner gaseous disc once it has settled into a relatively stable orientation."711 Hence the resulting distribution of stars is corresponcingly thin (although in S2 it is significantly thicker than 81. ef," Hence the resulting distribution of stars is correspondingly thin (although in S2 it is significantly thicker than S1, cf."712 Figure 6))., Figure \ref{fig:S1S2globus}) ).713 Such a distribution is broacly consistent with the observed. orbits of voung massive stars in the clockwise disc in the GC (7).., Such a distribution is broadly consistent with the observed orbits of young massive stars in the clockwise disc in the GC \citep{PaumardEtal06}.714 Due to gaseous disc midplane changes coupled with quicker fragmentation. then. the stellar disc in S4 is much thicker. with 7A-—I.," Due to gaseous disc midplane changes coupled with quicker fragmentation, then, the stellar disc in S4 is much thicker, with $H/R\sim 1$."715 This particular simulation thus fails to account for the most prominent. feature of the observational data: namely. a thin inner stellar disc.," This particular simulation thus fails to account for the most prominent feature of the observational data; namely, a thin inner stellar disc."716 Simulation S5 is somewhat distinct from all the rest due to its small impact. parameter. 6. between the gas clouds (see ‘Table 1).," Simulation S5 is somewhat distinct from all the rest due to its small impact parameter, $b$, between the gas clouds (see Table 1)."717 Xs a result of this fact and the relatively laree value of 3. a higher degree of mixing is achieved.," As a result of this fact and the relatively large value of $\beta$, a higher degree of mixing is achieved."718 Figure 12. shows a snapshot from this simulation. showing both gas and stars.," Figure \ref{fig:S5} shows a snapshot from this simulation, showing both gas and stars."719lensing. Sect.,"lensing, Sect."720 3. presents the adaptive contouring algorithm and Sect., \ref{sec:adaptive} presents the adaptive contouring algorithm and Sect.721 + shows how to determine the magnification of an extended source star from the image contour line.," \ref{sec:area}722 shows how to determine the magnification of an extended source star from the image contour line."723 Sect., Sect.724 5. discusses the critical curves and caustics of a binary point-mass lens in order to derive an algorithm for finding a point inside an image stretching over a critical curve., \ref{sec:binary} discusses the critical curves and caustics of a binary point-mass lens in order to derive an algorithm for finding a point inside an image stretching over a critical curve.725 Example light curves are shown in Sect., Example light curves are shown in Sect.726 6 before the paper concludes with a short summary in Sect. 7.., \ref{sec:lightcurves} before the paper concludes with a short summary in Sect. \ref{sec:summary}.727" Light received from a source object at distance Ds is bent due to the gravitational field of a thin sheet of matter at a distance 2i, with surface mass density X(£/) by the angle (2) For 8 denoting the true source position angle and @ the apparent position angle of its observed images. this implies theequation where w@=6/45. y=(DL/Ds) (9/0). and with (z)=ο0)2XGD0,2 )."," Light received from a source object at distance $D_\rmn{S}$ is bent due to the gravitational field of a thin sheet of matter at a distance $D_\rmn{L}$ with surface mass density $\Sigma({\vec \xi}')$ by the angle \citep{Schneider:theory}728 For $\vec \beta$ denoting the true source position angle and $\vec \theta$ the apparent position angle of its observed images, this implies the where $\vec x = \vec \theta/\theta_0$, $\vec y = (D_\rmn{L}/D_\rmn{S})\,(\vec \beta/\theta_0)$ , and with $\kappa({\vec x}') = (D_\rmn{L}\theta_0)^2\,\Sigma(D_\rmn{L} \theta_0729{\vec x}')$ ."730 It provides a surjective mapping of the image position a to the source position g. but the lack of injectivity means that a source may have more than just one image.," It provides a surjective mapping of the image position $\vec x$ to the source position $\vec y$, but the lack of injectivity means that a source may have more than just one image."731 At x. detined by two images merge. so that the number of images changes if and only if the source passes a pointy.=ylae).," At ${\vec x}_\rmn{c}$, defined by two images merge, so that the number of images changes if and only if the source passes a ${\vec y}_\rmn{c} = {\vec y}({\vec x}_\rmn{c})$."732 In general. critical points and caustic points form. closed curves. known as and. unless. these degenerate into a point in special cases.," In general, critical points and caustic points form closed curves, known as and unless these degenerate into a point in special cases."733 The conservation of surface brightness by gravitational lensing. It) Tyla]. /mplicsthattbetolalmagni ficationofapointsourecisgive so that it Uydiverges if the source comes to He on a caustic.," The conservation of surface brightness by gravitational lensing, $I(\vec x) = I[\vec y(\vec x)]$ , implies that the total magnification of a point source is given by so that it diverges if the source comes to lie on a caustic."734 For an extended source. as a consequence of Liouville's theorem. source and lens contours of same brightness correspond.," For an extended source, as a consequence of Liouville's theorem, source and lens contours of same brightness correspond."735 This means that if a source contour is described byan implicit function f(y:p)=0. with p being a parameter vector specifying the contour. then all its image contours are given by Py(a):p]= 0.," This means that if a source contour is described byan implicit function $F(\vec y; \vec p) = 0$, with $\vec p$ being a parameter vector specifying the contour, then all its image contours are given by $F[\vec y(\vec x); \vec p] = 0$ ."736 Therefore. a contour plot provides image contour lines without inversion of the lens equation (2)..," Therefore, a contour plot provides image contour lines without inversion of the lens equation \citep{SK87}."737 A search grid in the image plane needs to be large enough to cover all images and dense enough. so that no images or holes in the images are missed.," A search grid in the image plane needs to be large enough to cover all images and dense enough, so that no images or holes in the images are missed."738 While fulfilling these two condition. the grid resolution should be kept as low as possible.," While fulfilling these two condition, the grid resolution should be kept as low as possible."739 Obviouslv. an adaptive grid that just provides higher resolution in regions where this is required turns out to be superior to a fixed-resolution grid.," Obviously, an adaptive grid that just provides higher resolution in regions where this is required turns out to be superior to a fixed-resolution grid."740 Such an adaptive grid can be built by hierachically nesting squares that represent parts of the image plane., Such an adaptive grid can be built by hierachically nesting squares that represent parts of the image plane.741In order to approximate the position of the contour line. a ray is shot to the corresponding image position given by the detined mapping. Eq. (5).,"In order to approximate the position of the contour line, a ray is shot to the corresponding image position given by the defined mapping, Eq. \ref{eq:lenseq}) ),"742 and it is noted whether the corresponding true source position falls inside or outside the source contour., and it is noted whether the corresponding true source position falls inside or outside the source contour.743 In this respect. the proposed algorithm resembles the ray-shooting approach (22). but the indicator of whether the light ray hits the source is kept with the image position.," In this respect, the proposed algorithm resembles the ray-shooting approach \citep{KRS:rayshooting,SchneiWei:AGN}, but the indicator of whether the light ray hits the source is kept with the image position."744 In fact. with the conservation of surface brightness. as pointed out in Sect. 2..," In fact, with the conservation of surface brightness, as pointed out in Sect. \ref{sec:gravlens},"745 the image position has the nthjme relation (inside or outside) with respect to the image contour as the corresponding source position to the source contour., the image position has the same relation (inside or outside) with respect to the image contour as the corresponding source position to the source contour.746 A square with the inside/outside indicators for its corners constitutes the elementary datastructure. out of which all relevant information about the image plane is constructed.," A square with the inside/outside indicators for its corners constitutes the elementary datastructure, out of which all relevant information about the image plane is constructed."747 Being characterized by its four indicators. there are 24=16 types of elementary squares. shown in Fig. l..," Being characterized by its four indicators, there are $2^4 = 16$ types of elementary squares, shown in Fig. \ref{fig:square_type}."748 12 of these elementary squares. namely those of type “corner in/out’ or “half-in/half-out’. define part of the contour line. which crosses two edges of the square whose corners have different status.," 12 of these elementary squares, namely those of type 'corner in/out' or 'half-in/half-out', define part of the contour line, which crosses two edges of the square whose corners have different status."749 If two opposite corners are found to be inside and the other two corners outside. the contour needs to cross all four edges. but it is not clear which edges it connects.," If two opposite corners are found to be inside and the other two corners outside, the contour needs to cross all four edges, but it is not clear which edges it connects."750 Therefore. such indecisive squares need to be subdivided in order to improve the resolution (this is rule 1 below).," Therefore, such indecisive squares need to be subdivided in order to improve the resolution (this is rule 1 below)."751" The length of an edge of a square of depth & is 2)""a. where a denotes a unit size."," The length of an edge of a square of depth $k$ is $2^{-k} a$, where $a$ denotes a unit size."752" Squares are nested so that each square S$?nn of depth & either contains + subsquares of depth &|.1 each sharing a differentuc corner with ""mης or δη| does not have any subsquare."," Squares are nested so that each square $S^k_{mn}$ of depth $k$ either contains 4 subsquares of depth $k+1$ each sharing a different corner with $S^k_{mn}$, or $S^k_{mn}$ does not have any subsquare."753 The minimal and initial data structure consists of a square of depth 1 with 4 subsquaresofdepth 0., The minimal and initial data structure consists of a square of depth $-1$ with 4 subsquaresofdepth 0.754 The version of the adaptive contouring algorithm presented here uses just two operations on the data structure. and zxccnz. illustrated in Fig. 2.," The version of the adaptive contouring algorithm presented here uses just two operations on the data structure, and , illustrated in Fig. \ref{fig:operations}."755.While creates, .While creates756simulations provide the ideal tool to describe in detai how metals are produced. within galaxies and. distribute during the hierarchical assembly of a cluster.,simulations provide the ideal tool to describe in detail how metals are produced within galaxies and distributed during the hierarchical assembly of a cluster.757 Besides semianalytical approaches (e.g. De Lucia. Waullman White 2003). attempts to include. within hyerodynamica simulations. star formation. SN energy feedback and meta enrichment. from tvpela and LL SN. have been. pursue bv dillerent. authors (Aguirre ct al.," Besides semi–analytical approaches (e.g. De Lucia, Kauffman White 2003), attempts to include, within hydrodynamical simulations, star formation, SN energy feedback and metal enrichment from type–Ia and II SN, have been pursued by different authors (Aguirre et al."758 2001: Lia. Portinari Carraro 2002: Valdarnini 2002: Kawata Gibson 2003: Ixobavashi 2003: Fissera Scannapieco 2003).," 2001; Lia, Portinari Carraro 2002; Valdarnini 2002; Kawata Gibson 2003; Kobayashi 2003; Tissera Scannapieco 2003)."759" It is however clear that such approaches rely on the capability. of the numerical codes to provide a physically. sound. description of the relevant “subgrid"" processes.", It is however clear that such approaches rely on the capability of the numerical codes to provide a physically sound description of the relevant “sub–grid” processes.760 In thisLeffer we present the first. results from our hydrodynamical simulations of clusters based on the implementation of chemical enrichment in the GADGET code (Springel. Yoshida White 2001).," In this we present the first results from our hydrodynamical simulations of clusters based on the implementation of chemical enrichment in the GADGET code (Springel, Yoshida White 2001)."761 Our chemo-dvnamical version of GADGET combines the rather advanced: treatment of star formation and. SN. feedback. proposed by Springel Lernquist (2003a. 81103 hereafter). to a careful description of the role of tvpela ancl LL SN in releasing metalenriched gas into the diffuse medium.," Our chemo-dynamical version of GADGET combines the rather advanced treatment of star formation and SN feedback, proposed by Springel Hernquist (2003a, SH03 hereafter), to a careful description of the role of type–Ia and II SN in releasing metal–enriched gas into the diffuse medium."762 In the following. when expressing the ICM metal abundances in solar units. we assume the photospheric abundance provided by CGrevesse Sauval (1998).," In the following, when expressing the ICM metal abundances in solar units, we assume the photospheric abundance provided by Grevesse Sauval (1998)."763 Our simulations are based on an evolution of. (Springel ct al., Our simulations are based on an evolution of (Springel et al.764 2001). a parallel Tree|SPILL code with fully adaptive timestepping.," 2001), a parallel Tree+SPH code with fully adaptive time–stepping."765 As a starting point. we used a version. of GADGET. kindly provided by V. Springel. which includes an entropyconserving integration scheme. radiative cooling. the effect of a uniform and evolving UV background. (Llaardt Aladau 1909). star formation from a multiphase interstellar medium. and a prescription for galactic winds triggered by SN explosions (sec SIIO3 for à detailed description).," As a starting point, we used a version of GADGET, kindly provided by V. Springel, which includes an entropy–conserving integration scheme, radiative cooling, the effect of a uniform and evolving UV background (Haardt Madau 1999), star formation from a multiphase interstellar medium and a prescription for galactic winds triggered by SN explosions (see SH03 for a detailed description)."766 In the original version of the code. the energev release and a global metallicity was produced. only by SNIL under the instantanousrecvcling. approximation (LRA).," In the original version of the code, the energy release and a global metallicity was produced only by SNII under the instantanous–recycling approximation (IRA)."767 The GADGET code has been suitably mocified. so as to correctly include the lifetimes of dillerent. stellar populations. to follow metal production from both SNlIa and Η. while self.consistently introducing the dependence of the cooling function on metallicitv.," The GADGET code has been suitably modified, so as to correctly include the life–times of different stellar populations, to follow metal production from both SNIa and II, while self–consistently introducing the dependence of the cooling function on metallicity."768 A detailed description of the implementation of these algorithms will be presented in a forthcoming paper Cl'ornatore ct al., A detailed description of the implementation of these algorithms will be presented in a forthcoming paper (Tornatore et al.769 2004. in. preparation). while we provide here à short descriptions of the most relevant. features of the code.," 2004, in preparation), while we provide here a short descriptions of the most relevant features of the code."770 In order to maintain the general approach of the multiphase model by SLIIQ3. we still treat under the LRA stars with masses ον20M.. while accounting for the different life-times of stars of smaller mass (Matteucci DPadovani 1993).," In order to maintain the general approach of the multiphase model by SH03, we still treat under the IRA stars with masses $>20\,M_\odot$, while accounting for the different life-times of stars of smaller mass (Matteucci Padovani 1993)."771 Within the stochastic approach to star formation (SIIQ3). each star particle is considered as a single stellar population (SSP).," Within the stochastic approach to star formation (SH03), each star particle is considered as a single stellar population (SSP)."772 For each SSP we compute the number of stars turning into SNIL and la at each time-step after its creation., For each SSP we compute the number of stars turning into SNII and Ia at each time-step after its creation.773 The SNla are associated to. binary svstems whose components are in the SAZ. mass range (Ciregeio venzini 1983). while SNIL arise from stars with mass SAL. (ef.," The SNIa are associated to binary systems whose components are in the $8\,M_\odot$ mass range (Greggio Renzini 1983), while SNII arise from stars with mass $>8\,M_\odot$ (cf."774 also Lia et al., also Lia et al.775 2002. who adopt the lower mass hreshold of GAL. for SNIL).," 2002, who adopt the lower mass threshold of $6M_{\odot}$ for SNII)."776 Besides SNe. which release energy and metals. we also account for planctary nebulae (PN).," Besides SNe, which release energy and metals, we also account for planetary nebulae (PN)."777 They contribute to metal production. but not to the energv feedback. and are identified with those stars. not urning into SNla. in the mass range SM..," They contribute to metal production, but not to the energy feedback, and are identified with those stars, not turning into SNIa, in the mass range $\,M_\odot$."778 We use the analvtical fitting formulas for stellar vields of SNla. SNIL ancl PNe as provided by Reechi et al. (," We use the analytical fitting formulas for stellar yields of SNIa, SNII and PNe as provided by Recchi et al. ("7792001). ancl based on he original nucleosvnthesis computations of Nomoto et al. (,"2001), and based on the original nucleosynthesis computations of Nomoto et al. ("7801997. using their W model). Woosley Weaver (1995) and Benzin Voli (1981).,"1997, using their W7 model), Woosley Weaver (1995) and Renzini Voli (1981)."781 The ormulation for the SNla rate has been calculated. as in Alatteuccei Reeehi (2001)., The formulation for the SNIa rate has been calculated as in Matteucci Recchi (2001).782 Besides HE and Ho. the current version of the code follows the production. of Fe. ο. ο Si Mg. S. and can be casily modified. to include other metal species.," Besides H and He, the current version of the code follows the production of Fe, O, C, Si, Mg, S, and can be easily modified to include other metal species."783 Once. produced. by à star. particle. metals are spread over the same number of neighbours. 82. used [or the SPLL implementation. also using the same kernel.," Once produced by a star particle, metals are spread over the same number of neighbours, 32, used for the SPH implementation, also using the same kernel."784" In this wav. we find that 90 per cent of the metals are distributedDl withinM a gas mass of. 5.4.1075.""ΑΙ..."," In this way, we find that 90 per cent of the metals are distributed within a gas mass of $5.4\times 10^9785h^{-1}M_\odot$."786 We. have verified that using a twice as Large number of neighbors to spread metals results in a twice as large gas mass for metal mixing. while final results on the amount and distribution of metals (see below) are left almost unchanged.," We have verified that using a twice as large number of neighbors to spread metals results in a twice as large gas mass for metal mixing, while final results on the amount and distribution of metals (see below) are left almost unchanged."787 As for the energy release. each SN is assumed to produce 10ergs.," As for the energy release, each SN is assumed to produce $10^{51}$ ergs."788 Insteacl of assuming any specific value for the thermalization ellicieney. of the energy. released by SN. we prefer to dump all the energy to the surrounding σας particles and. leave to the simulation the computation of the radiation losses.," Instead of assuming any specific value for the thermalization efficiency of the energy released by SN, we prefer to dump all the energy to the surrounding gas particles and leave to the simulation the computation of the radiation losses."789 Since the physical processes determining the actual SN Mlicieney are below the resolution scale of our simulations. 16 rationale behind our choice is to leave to the sub.grid multiphase niocel by 1109 establishing how much of this energy enters in regulating the star formation process.," Since the physical processes determining the actual SN efficiency are below the resolution scale of our simulations, the rationale behind our choice is to leave to the sub–grid multiphase model by SH03 establishing how much of this energy enters in regulating the star formation process."790 We normalize the IMEs in the mass range 0.1.LOO A7..., We normalize the IMFs in the mass range 0.1–100 $M_\odot$.791 Owing to the uncertainty in modelling vields for very massive stars. we take vields to be independent. of mass above 40M...," Owing to the uncertainty in modelling yields for very massive stars, we take yields to be independent of mass above $40\,M_\odot$."792 While any uncertainty in the vields of such massive stars has a negligible effect for a Salpeter IML (Salpeter 1955. 855 hereafter). their accurate description (e.g. Phiclemann et al.," While any uncertainty in the yields of such massive stars has a negligible effect for a Salpeter IMF (Salpeter 1955, S55 hereafter), their accurate description (e.g. Thielemann et al."793 1996: Leger Woosley 2002) is required. when using a topheavier EME., 1996; Heger Woosley 2002) is required when using a top–heavier IMF.794 We note that our scheme to distribute metals in the ICM does not include the effect. o£ dillusion., We note that our scheme to distribute metals in the ICM does not include the effect of diffusion.795 Lia et al. (, Lia et al. (7962002) included the ellect. o£ cülfusion driven by SN blast waves (see Thornton et al.,2002) included the effect of diffusion driven by SN blast waves (see Thornton et al.797 1998) in their SPILL simulations with chemical enrichment., 1998) in their SPH simulations with chemical enrichment.798 Although this cllect is quite important to describe the dilfusion of metals within the interstellar medium. it is likely to play a minor role on . . ≱∖⋯↓⋖⋅≱∖⋜↧∣⋯∖⇁∢⋅↥⇂↥⋖⋅↿∙∖⇁↓≻⊔∼⋜↧↓↓⋅∢⊾⊳∖∪↓⋯↓∪⊔⊳∖⋯↓⋖⋅⊳∿↓∪∣∣↳↓≻≼⋱∪⊓↓∐⊾ 1 ⋅ cluster simulations that we are discussing here.," Although this effect is quite important to describe the diffusion of metals within the interstellar medium, it is likely to play a minor role on scales above the typical resolution scale, $\sim 10\,h^{-1}$ kpc, of the cluster simulations that we are discussing here."799zero vields negligibly dillerent results.,zero yields negligibly different results.800" For our final solution (solution 19 of Table 5) we adopt. 2,=8.5kpe (see below) and corresponding values of O7 and OF from fig 5 of PMD together with the values of i, and ον, from PMD.", For our final solution (solution 19 of Table 5) we adopt $R_{\rm o}=8.5\ {\rm kpc}$ (see below) and corresponding values of $\Theta_{\rm o}''$ and $\Theta_{\rm o}'''$ from fig 5 of PMB together with the values of $u_{\rm o}$ and $v_{\rm o}$ from PMB.801" Thus we adopt from the proper JA-—1482cz054kmsl. B=1237EO61kms3 and £2,=21.1940.87kms These results are. essentially independent of the acloptecl distance scale or of scaling the PAIB values of R.A"" and Q7 (compare solutions 15. 17 and 19 of table 5)."," Thus we adopt from the proper $ A = 14.82 \pm 0.84\ \rm km\ s^{-1}$, $ B = -12.37 \pm 0.64\ \rm km\ s^{-1}$ and $\Omega _{\rm o} = 27.19 \pm 0.87\ \rm km\ s^{-1}$ These results are essentially independent of the adopted distance scale or of scaling the PMB values of $R_{\rm o}, \Theta''_{\rm o}$ and $\Theta'''_{\rm o}$ (compare solutions 15, 17 and 19 of table 5)."802 Comparison of this adopted value of; with the value from racial velocities (15.92 £0.34) leacs (via fig 5 of PAIB), Comparison of this adopted value of $A$ with the value from radial velocities $15.92 \pm 0.34$ ) leads (via fig 5 of PMB)803Vel 668826) was classified as BpSi bv. ancl was reported. as an eclipsing binary by (1995).,Vel 68826) was classified as BpSi by and was reported as an eclipsing binary by .804.. In fact it is one of the only two double-Hined eclipsing binaries with a Bpsi component known to date 2004)., In fact it is one of the only two double-lined eclipsing binaries with a BpSi component known to date .805. From light-time effect on the times of minima. deduced the presence of a third body.," From light-time effect on the times of minima, deduced the presence of a third body."806 In our previous paper 2006).. using FEROS spectroscopic observations. we discovered that this system is actually a spectroscopic quadruple system: with components close to the ZAXMS.," In our previous paper , using FEROS spectroscopic observations, we discovered that this system is actually a spectroscopic quadruple system with components close to the ZAMS."807 The four stars form. two close spectroscopic pairs (periods of 1.58 and 4.15 days) bound gravitationally to each other in a wide eccentric orbit with a period of vr., The four stars form two close spectroscopic pairs (periods of 1.58 and 4.15 days) bound gravitationally to each other in a wide eccentric orbit with a period of 41 yr.808 In that paper we combined our radial velocity (RV) measurements41 with the available photometric data to derive orbital parameters for both binary systems and to calculate the absolute parameters of the eclipsing svstem., In that paper we combined our radial velocity ) measurements with the available photometric data to derive orbital parameters for both binary systems and to calculate the absolute parameters of the eclipsing system.809 For the first time. direct determination of the radius and the mass was obtained for a Όροι star.," For the first time, direct determination of the radius and the mass was obtained for a BpSi star."810 In this work we present high-resolution. high signal-to-noise UVIES spectra. which were used to perform. an abundance analysis of all four components of this multiple system.," In this work we present high-resolution, high signal-to-noise UVES spectra, which were used to perform an abundance analysis of all four components of this multiple system."811 In 22 we present the observations and describe the reconstruction of the component spectra., In 2 we present the observations and describe the reconstruction of the component spectra.812 In 33 we analyze the spectral characteristics of cach component and present the results of the abundance analysis., In 3 we analyze the spectral characteristics of each component and present the results of the abundance analysis.813 In the last Section we discuss the main results and the occurrence of chemical peculiarities in binary ancl multiple svstenis., In the last Section we discuss the main results and the occurrence of chemical peculiarities in binary and multiple systems.814 Three spectra were obtained in service mode with UVES at VLI-UT2 telescope in October 2005., Three spectra were obtained in service mode with UVES at VLT-UT2 telescope in October 2005.815 The spectra have been taken with the aarcsec slit for the blue arm and the aarcsec slit for the red arm on three consecutive nights with two cdillerent clichroics to achieve the highest UVES resolution of 110.000 in the red spectral region ancl 50.000 in the blue spectral region.," The spectra have been taken with the arcsec slit for the blue arm and the arcsec slit for the red arm on three consecutive nights with two different dichroics to achieve the highest UVES resolution of 110,000 in the red spectral region and 80,000 in the blue spectral region."816 We used exposure times of 20.30 min. obtaining à S/N ratio above 200 in the spectral rangeSOOOA.," We used exposure times of 20–30 min, obtaining a S/N ratio above 200 in the spectral range."817. These spectra are analyzed here along with five FEROS spectra described in the previous paper bv(2006)., These spectra are analyzed here along with five FEROS spectra described in the previous paper by.818. To calculate. separate spectra for the four components of the system and to measure their RVs. the iterative method. described by was adapted for the multiple system AO Vel.," To calculate separate spectra for the four components of the system and to measure their RVs, the iterative method described by was adapted for the multiple system AO Vel."819 This algorithm computes the spectra of the individual components and the RVs iteratively., This algorithm computes the spectra of the individual components and the RVs iteratively.820 In each step the computed spectra are used to remove the spectral features of all but one component from the observed spectra., In each step the computed spectra are used to remove the spectral features of all but one component from the observed spectra.821 Ehe resulting single-lined spectra, The resulting single-lined spectra822secondary pair production im pulsars requires a source of photons will energies 2l1 MeV. The sources considered in conventional models are curvature enüssion and resonant Thomson scattering by primary. particles.,"Secondary pair production in pulsars requires a source of photons with energies $>1\,$ MeV. The sources considered in conventional models are curvature emission and resonant Thomson scattering by primary particles."823 In an oscillating model LAE is an additional possibility., In an oscillating model LAE is an additional possibility.824 For LAE to be viable as the source of secondary pairs. (wo conditions need {ο be satisfied: the photon energv must exceed an MeV. and the power in LAE must be sufficient to account for the required number of pairs.," For LAE to be viable as the source of secondary pairs, two conditions need to be satisfied: the photon energy must exceed an MeV, and the power in LAE must be sufficient to account for the required number of pairs."825 Consider a model in which there is a large number of localized. transient LAEWs in the polar cap region. with the pairs in the LAEW created through LAE.," Consider a model in which there is a large number of localized, transient LAEWs in the polar cap region, with the pairs in the LAEW created through LAE."826" Let the number density of pairs be a multiplicity. AZ. times the Goldreich-Julian density. so that the lrequency of the LAEW is eascCUO,0,)!s where Q,=2z/P is (he rotation Irequency. of a pulsar with period P. and Q,=(nc?/h)uB/D.) is the evelotron frequency. with B,=44x10""T the Schwinger field."," Let the number density of pairs be a multiplicity, $M$, times the Goldreich-Julian density, so that the frequency of the LAEW is $\omega_{\rm max}\sim (M\Omega_r\Omega_c)^{1/2}\gamma_{\rm max}^{3/2}$, where $\Omega_r=2\pi/P$ is the rotation frequency of a pulsar with period $P$, and $\Omega_c=(mc^2/\hbar)(B/B_c)$ is the cyclotron frequency, with $B_c=4.4\times10^9\rm\,T$ the Schwinger field."827 The threshold condition. «y27/2mech. requires110. where P is in seconds.," The threshold condition, $\omega_{\rm max}>2mc^2/\hbar$, requires, where $P$ is in seconds."828 The fraction of the οποιον lost bv a LAEW as it propagates outward through the pulsar magnetosphere can be estimated by mulliplving this damping decrement by the number of oscillations before the LAEW leaves the magnetosphere., The fraction of the energy lost by a LAEW as it propagates outward through the pulsar magnetosphere can be estimated by multiplying this damping decrement by the number of oscillations before the LAEW leaves the magnetosphere.829" Assuming propagation al close to the speed of light this number is of order Q/Q,.", Assuming propagation at close to the speed of light this number is of order $\Omega/\Omega_r$.830" Hence. the fraction of the energv lost to LAE is of order (r5?/O,c)suusoALryQ,/¢=MalB/B.). where a221/137 is the line structure constant."," Hence, the fraction of the energy lost to LAE is of order $(r_0\Omega^2/\Omega_rc)\gamma_{\rm max}\sim Mr_0\Omega_c/c=M\alpha(B/B_c)$, where $\alpha\approx1/137$ is the fine structure constant."831 We conclude that LAE is energetically, We conclude that LAE is energetically832optical outburst is difficult to assess.,optical outburst is difficult to assess.833 While the optical outburst is naturally interpreted as an accretion outburst (with the optical luminosity due to increased luminosity of the accretion shock). why this should cause a hot plasma component to appear in the star is not clear.," While the optical outburst is naturally interpreted as an accretion outburst (with the optical luminosity due to increased luminosity of the accretion shock), why this should cause a hot plasma component to appear in the star is not clear."834 Indeed. accreting YSOs in the COUP sample are statistically less X-ray luminous than non-accreting ones (Preibischetal.. 2005)) and less prone to show flares from large magnetic structures (Favataetal.. 2005)).," Indeed, accreting YSOs in the COUP sample are statistically less X-ray luminous than non-accreting ones \citealp{pkf+2005}) ) and less prone to show flares from large magnetic structures \citealp{ffr+2005}) )."835We parameterise the global densities in which individual objects are embedded by measuring their distance to notorious landmarks in the simulation.,We parameterise the global densities in which individual objects are embedded by measuring their distance to notorious landmarks in the simulation.836" Possible choices For lucdanarks inelucle galas, clusters aud voids. which correspond to rare fluetuations in the density Ποια."," Possible choices for landmarks include galaxy clusters and voids, which correspond to rare fluctuations in the density field."837 Pilaments could also be selected as Iandinarks bu in these structures are located at the void walls in our niunerieal principlesimulation (Le. 0.8 ont) and we therelore do no use them in our analysis., Filaments could also be selected as landmarks but in principle these structures are located at the void walls in our numerical simulation (i.e. $0.8-1.2 r_{void}$ ) and we therefore do not use them in our analysis.838 In order to illustrate the meaning of a global density parametrised this way. notice that galaxies locatec on aspherical shell centred in. lor instanee. a Cluster al galaxies could in principle be embedded in a wide range of local densities. depending on whether they are field galaxies or part of filaments or groups.," In order to illustrate the meaning of a global density parametrised this way, notice that galaxies located on a spherical shell centred in, for instance, a Cluster of galaxies could in principle be embedded in a wide range of local densities, depending on whether they are field galaxies or part of filaments or groups."839 Lor a fixed global environment the local environmen can ciller significantlv., For a fixed global environment the local environment can differ significantly.840 DAL haloes in the numerical simulation follow quite approximately NEW profiles. (Navarro.Prenk&White1996).. where haloes of similar concentrations ean be sealed to a single prolile Using a scale radius. reog. which eneloses an overdensity of 200 times the critical density in the aniverse.," DM haloes in the numerical simulation follow quite approximately NFW profiles \citep{nfw}, where haloes of similar concentrations can be scaled to a single profile using a scale radius, $r_{200}$, which encloses an overdensity of $200$ times the critical density in the universe."841 Once this scaling is applied. the spherical shells distant by 7/rooo Irom any halo centre are characterised) by similar overdensities.," Once this scaling is applied, the spherical shells distant by $r/r_{200}$ from any halo centre are characterised by similar overdensities."842 Phis approximation js also valid lor the full population of haloes whieh presents a Barrow range ol possible concentrations (σοςetal.1998]., This approximation is also valid for the full population of haloes which presents a narrow range of possible concentrations \citep{hus}.843 As the latter becomes an even better approximation when the population of haloes is restricted to a narrow. range al masses. we select as Iuncdanarks lor global density estimators haloes with Abs olehTAL.. for a total of T0 selected DAL haloes.," As the latter becomes an even better approximation when the population of haloes is restricted to a narrow range of masses, we select as landmarks for global density estimators haloes with $M>10^{13}$ $^{-1}M_{\odot}$, for a total of $70$ selected DM haloes."844 We then proceed ta label galaxies according to their distance. in terms of rogo. to the closest DM halo within this ," We then proceed to label galaxies according to their distance, in terms of $r_{200}$, to the closest DM halo within this sample."845"From now on. we divide galaxies in. 1 stthsamples at different. sample.distances from halo centres. which we will reler to as Ay, to ga (with limiting values at rfrouy=0.0.1.5.5.9 and 20)."," From now on, we divide galaxies in $4$ subsamples at different distances from halo centres, which we will refer to as $R_{H1}$ to $R_{H4}$ (with limiting values at $r/r_{200}=0.0,1.5,5,9$ and $20$ )."846 The corresponding average DAL density around galaxies in cach stthbsample ranges Irom σε200;%¢« lo 7per where pec ds the critical mass density.," The corresponding average DM density around galaxies in each subsample ranges from $\approx 200 \rho_C$ to $\approx \rho_C$, where $\rho_C$ is the critical mass density."847 A similar principle applies to voids where their density proliles can be scaled using the void radius ρα (Padilla.Ceeca-relli&Lambas 2005): the profiles approach the average densit in the Universe at ορα£1.5 (Patirietal.2006)..," A similar principle applies to voids where their density profiles can be scaled using the void radius $r_{void}$ \citep{pad}; the profiles approach the average density in the Universe at $r/r_{void} 848\approx 1.5$ \citep{pat}."849 The void. identification algorithin we adopt. corresponds to the one deseribed tn Padilla.Ceccarelli&Lambas(2005).. and consists of a search of underdense spheres of varving radii within the periodic simulation box. satisbving 0=pop0.9.," The void identification algorithm we adopt corresponds to the one described in \citet{pad}, and consists of a search of underdense spheres of varying radii within the periodic simulation box, satisfying $\delta=\frac{\rho-\left<\rho\right>}{\left<\rho\right>}<-0.9$."850" In the simulation x woe identify a total of TO voids"" with 7,54 lh 4Alpe. each containing in average a total of ©190 galaxies in the range ρω5USN1.2."," In the simulation box we identify a total of $70$ voids with $r_{void}>4$ $^{-1}$ Mpc, each containing in average a total of $\simeq 180$ galaxies in the range $r/r_{void}=0.8-1.2$."851 We use these voids to make a second xuwaneterisation of global densities lor the semi-analytic galaxies using Pérou., We use these voids to make a second parameterisation of global densities for the semi-analytic galaxies using $r/r_{void}$.852" We define 1 distance ranges. relerred to as a4 to Haas delimitec by the values οδρ=0.0.55.0.85. 1.05 and 1,1: he average overdensity in these samples ranges [rom àΟσο: WFpet."," We define $4$ distance ranges, referred to as $R_{V1}$ to $R_{V4}$, delimited by the values $r/r_{void}=0,0.55,0.85$, $1.05$ and $1.4$; the average overdensity in these samples ranges from $\approx 0.05 \rho_C$ to $\approx \rho_C$."853 The resulting distributions of normalised distances to lialocs and voids lor the seimi-analvtie galaxies in the simulation are shown in Figure 1z the vertical long-dashed lines indicate the limits between clifferent global density samples selected according to the distance to haloes (left. panel) ancl voids (right) in the simulation., The resulting distributions of normalised distances to haloes and voids for the semi-analytic galaxies in the simulation are shown in Figure \ref{fig:fig1}; the vertical long-dashed lines indicate the limits between different global density samples selected according to the distance to haloes (left panel) and voids (right) in the simulation.854 Solid lines show the results lor the full sample of galaxies in the simulation. dotted lines show central galaxies (notice the lack of near the halo centres. indicating the minimum distance objectsbetween haloes in the simulation). and typicaldashed lines to satellite galaxies.," Solid lines show the results for the full sample of galaxies in the simulation, dotted lines show central galaxies (notice the lack of objects near the halo centres, indicating the typical minimum distance between haloes in the simulation), and dashed lines to satellite galaxies."855 Except for the lack of central galaxies near halo. centres. the of these distributions do not change significantly with the shapesgalaxy tape.," Except for the lack of central galaxies near halo centres, the shapes of these distributions do not change significantly with the galaxy type."856 ‘Lhe problem of celining a measure of the local density. aro galaxies has multiple possible solutions. depending on the relevan quantities that are to be associated to this estimate.," The problem of defining a measure of the local density around galaxies has multiple possible solutions, depending on the relevant quantities that are to be associated to this estimate."857 On the one mand the standard approach of estimating a density using a fixec voltune around a galaxy eimsures a fixed seale. but the extremelh wide dynamic range of densities (228 orders of magnitude) has he drawback of procucing low signal to noise estimates [or ow densities. (dominated. by Poisson noise). ancl oversample meastrements at high cleusity values;," On the one hand the standard approach of estimating a density using a fixed volume around a galaxy ensures a fixed scale, but the extremely wide dynamic range of densities $\approx 8$ orders of magnitude) has the drawback of producing low signal to noise estimates for low densities (dominated by Poisson noise), and oversampled measurements at high density values."858 Several works apply this method either by using a gaussian kernel to smooth the density distribution (e.g. Balogh et al.," Several works apply this method either by using a gaussian kernel to smooth the density distribution (e.g. Balogh et al.,"859 2001a). or a more simple top-jx kernel of fixed. size (C'eccarelli:Padilla&Lambas20058).," 2004a), or a more simple top-hat kernel of fixed size \citep{cec}."860. A thoiee of local density estimate associated: to galaxv-galaxy interactions is the one taking into account the closest. neeghbors oL a galaxy by using the distance to the IN/ nearest neighbor., A choice of local density estimate associated to galaxy-galaxy interactions is the one taking into account the closest neighbors of a galaxy by using the distance to the $N^{th}$ nearest neighbor.861 The advantages from using such an estimator relies in the likely relation between galaxy interactions anc the SI in galaxies., The advantages from using such an estimator relies in the likely relation between galaxy interactions and the SF in galaxies.862" Mos observational studies adopting this latter approach use adaptive projected 2D density estimators such as Ss: lor instance. Baloghetal.(2001a) compute Xz using the distance to the filth neares neighbour brighter than AZ,=20 confined to a recdshilt slice of της+ to avoid biases Irom the linger-ol-eud elect."," Most observational studies adopting this latter approach use adaptive projected 2D density estimators such as $\Sigma_5$; for instance, \citet{bal} compute $\Sigma_5$ using the distance to the fifth nearest neighbour brighter than $M_r=-20$ confined to a redshift slice of $\pm 1000km s^{-1}$ to avoid biases from the finger-of-god effect."863" ""he alternatives to these approaches. generally applied to nammerieal simulations with Full 3-dimensional information. σασί on using an smoothing length proportional to the local particle oradaptive galaxy separation. or the galaxy-galasy distance (lor Smoothed Particle Lyclrodynamies. SPI. or adaptive loca density estimates given by Voronoi Pessellations (VI. Voronoi. 1908)."," The alternatives to these approaches, generally applied to numerical simulations with full 3-dimensional information, consist on using an adaptive smoothing length proportional to the local particle or galaxy separation, or the galaxy-galaxy distance (for Smoothed Particle Hydrodynamics, SPH), or adaptive local density estimates given by Voronoi Tessellations (VT, Voronoi, 1908)."864 With VI. each particle is associated to a domain volume πο that every point inside this volume is closer to the particle ad its centre than to any other particle: smoothing these density estimates with those of their immediate neighbors ean be used to obtain a reliable measure of the local density.," With VT, each particle is associated to a domain volume so that every point inside this volume is closer to the particle at its centre than to any other particle; smoothing these density estimates with those of their immediate neighbors can be used to obtain a reliable measure of the local density."865 This. particular measurement method. called the Literpolated Voronoi Density VD). can be very useful Lor identification af ος objects Platen.vandeWeygaert&Jones2007:Aragon-Calvuetal.Nevrinck.Cinedin&{αμα2005:CionzálezTheuns2009). in ..PI simulations and bas beenο shown to have a better resolution jur any other adaptive density estimate including SPIL kernel smoothing techniques (Schaap&vandeWergacrt2000:Pelu-oessyCtal.," This particular measurement method called the Interpolated Voronoi Density (IVD), can be very useful for identification of bound objects \citep{platen,aragon,ney,rob}866 in SPH simulations and has been shown to have a better resolution than any other adaptive density estimate including SPH kernel smoothing techniques \citep{schaap,pel}."867 2003). In the remainder of this work we adopt IVDs or our estimates of local densities in the numerical simulation., In the remainder of this work we adopt IVDs for our estimates of local densities in the numerical simulation.868 The left. panel o£ Figure 2. shows the IVD distributions of galaxies ott to 207200 Lrom the cluster centres (same as in the left xuiel ofFigure 1: the right panel shows the IVD distributions of galaxies ott to ἘνCeh from the void centres (as in the right xuiel of 1)., The left panel of Figure \ref{fig:fig2} shows the IVD distributions of galaxies out to $20r_{200}$ from the cluster centres (same as in the left panel of 1); the right panel shows the IVD distributions of galaxies out to $1.4r_{void}$ from the void centres (as in the right panel of 1).869 RegardlessΌσο of whether the haloes elt panel) or voids (right panel) are used as global density laucimarks. the distributions of local densities are similar since both selection criteria cover a large fraction ofthe volume ofthe simulation.," Regardless of whether the haloes (left panel) or voids (right panel) are used as global density landmarks, the distributions of local densities are similar since both selection criteria cover a large fraction of the volume of the simulation."870 Phe VD distributions show a clear bimodal behaviour. where the peak at hieh densities is dominated by satellite galaxies in massive DN 1alocs (long-dashed lines). and the density distribution of centra galaxies (short-dashecl) retleets the density field around their lios raloes. characterised by masses AfL04th 1AZ..," The IVD distributions show a clear bimodal behaviour, where the peak at high densities is dominated by satellite galaxies in massive DM haloes (long-dashed lines), and the density distribution of central galaxies (short-dashed) reflects the density field around their host haloes, characterised by masses $M\gtrsim 10^{11}$ $^{-1}M_{\odot}$."871 We use these density estimates tuo characterise the loca environment of the semi-analytie galaxies., We use these density estimates to characterise the local environment of the semi-analytic galaxies.872" La ραοπο, we wil separate galaxies in three local density bins. to be relerred to as the prow. pup and µε samples."," In particular, we will separate galaxies in three local density bins, to be referred to as the $\rho_{LOW}$ $\rho_{MID}$ and $\rho_{HIGH}$ samples."873 Lhe first tentative cuts in stellar densities are applied at 3.16x1027h. FAL. ? ane 6:1xLlotth ΑΗ selected: so as to lave a significan number of galaxies Alpein cach subsample.," The first tentative cuts in stellar densities are applied at $3.16 \times 10^{10}$ $^{-2}M_{\odot} $ $^{-3}$ and $6.31 \times 10^{11}$ $^{-2}M_{\odot} $ $^{-3}$, selected so as to have a significant number of galaxies in each subsample."874 Further restrictions in density may be needed in order to ensure a constant median IND in each sample studied., Further restrictions in density may be needed in order to ensure a constant median IVD in each sample studied.875" Por reference.the barvon density in the simulation is py=f,x2.810h. 72A. 7. with à barvon fraction fj= 0.037."," For reference,the baryon density in the simulation is $\rho_b = f_b \times 2.8 \times 10^{11}$ $^{-2}M_{\odot} $ $^{-3}$ , with a baryon fraction $f_b=0.037$ ."876"the volume of à single cell and NV,=398 is the number of ealaxies in the simulation box.",the volume of a single cell and $N_g=398$ is the number of galaxies in the simulation box.877 To allow a direct comparison of the simulation results with the observational data we construct synthetic Lva spectra as described in Bruscoliat... by tracing random LOS through the box.," To allow a direct comparison of the simulation results with the observational data we construct synthetic $\alpha$ spectra as described in Bruscoli, by tracing random LOS through the box."878 We then estimate the mean Lya transmitted [Dux as a function of the impact parameter Ar. computing the average on all the pixels at a distance from a galaxy in a given interval centered on Ar.," We then estimate the mean $\alpha$ transmitted flux as a function of the impact parameter $\Delta r$, computing the average on all the pixels at a distance from a galaxy in a given interval centered on $\Delta r$."879 The results are plotted in Fig.3 ancl compared with the data. represented by the black points.," The results are plotted in Fig.3 and compared with the data, represented by the black points."880 The three panels show the mean Lye lux computed on different samples of galaxies: all the 398 galaxies (left panel). ealaxies with mass above 2107 M. (central panel) and those with mass below 9.3107 M. (right panel).," The three panels show the mean $\alpha$ flux computed on different samples of galaxies: all the 398 galaxies (left panel), galaxies with mass above $2 \times 10^{10}$ $_\odot$ (central panel) and those with mass below $9.3 \times 10^8$ $_\odot$ (right panel)."881 Solid. dotted and. dotted-dashed lines are derived from the outputs of run A. run D and from a simulation analogous to run B. but with SERs boosted.," Solid, dotted and dotted-dashed lines are derived from the outputs of run A, run B and from a simulation analogous to run B, but with SFRs boosted."882hoc. Phe dillerence between run A and run BD in the left panel is marginal: we conclude that the mean cllect of the local photoionization on the Lya transmitted [lux is negligible when all galaxies are included in the analysis., The difference between run A and run B in the left panel is marginal: we conclude that the mean effect of the local photoionization on the $\alpha$ transmitted flux is negligible when all galaxies are included in the analysis.883 Massive galaxies are the best candidates for LBCs due to their high luminosity and clustering properties., Massive galaxies are the best candidates for LBGs due to their high luminosity and clustering properties.884 The SET. derived for the most massive galaxies in the MSPL simulation are in the range LO30 M.vr the highes values in the simulation.," The SFR derived for the most massive galaxies in the MSPH simulation are in the range $10\div30$ $_{\odot} {\rm yr}^{-1}$, the highest values in the simulation."885 Due to their high. luminosity. one would expect them to produce a strong impact on the ionization of the surrounding σας.," Due to their high luminosity, one would expect them to produce a strong impact on the ionization of the surrounding gas."886 However this is no the case., However this is not the case.887 Looking at the solid line in the central panel. obtained by neglecting the local emission. we note tha the hieh density characterizing the environment of massive galaxies suppresses the mean transmitted Dux with respec to the mean trend obtained. for all the galaxies in. the simulation (solid line in the left panel).," Looking at the solid line in the central panel, obtained by neglecting the local emission, we note that the high density characterizing the environment of massive galaxies suppresses the mean transmitted flux with respect to the mean trend obtained for all the galaxies in the simulation (solid line in the left panel)."888 The UV radiation emitted by these galaxies. using the SER. derived from the MSPL simulation (dotted. line in central panel). is not strong enough to enhance significantly the transparency. of he surrounding ICM. because of the high recombination rale in the denser environment.," The UV radiation emitted by these galaxies, using the SFR derived from the MSPH simulation (dotted line in central panel), is not strong enough to enhance significantly the transparency of the surrounding IGM, because of the high recombination rate in the denser environment."889 Higher values of SER. are necessary to produce a significant increase of the transmitted lux nearby massive galaxies., Higher values of SFR are necessary to produce a significant increase of the transmitted flux nearby massive galaxies.890 This can be seen from the x)osted. case. with SER in the range 100300 M. +.," This can be seen from the boosted case, with SFR in the range $100 \div 300$ $_\odot$ $^{-1}$."891 Although on the high side. the values are still plausible according to observations (Shapley.. 2001).," Although on the high side, the values are still plausible according to observations (Shapley, 2001)."892 Despite he sharp increase in the mean [να transmitted. (lux. vet his simulation does not match the AOS data.," Despite the sharp increase in the mean $\alpha$ transmitted flux, yet this simulation does not match the A03 data."893 This result sugeests that the mean transmissivity strongly depends on he galactic enviroment., This result suggests that the mean transmissivity strongly depends on the galactic enviroment.894 The right panel shows the trend for he Lya transmitted [lax obtained. selecting galaxies with mass <=9.8.107 M.., The right panel shows the trend for the $\alpha$ transmitted flux obtained selecting galaxies with mass $ \le 9.3 \times 10^8$ $_\odot$.895 In this case the transmissivity of he gas at Ar<O.7h+ Alpe has an opposite trend. with respect to the mean (left panel). ancl seems to follow the one observed by AOS.," In this case the transmissivity of the gas at $\Delta r < 0.7 h^{-1}$ Mpc has an opposite trend with respect to the mean (left panel), and seems to follow the one observed by A03."896 The increasing Dux at smaller distance rom the galaxy rellects the lower density in these regions., The increasing flux at smaller distance from the galaxy reflects the lower density in these regions.897 ere the UVB flux can counteract the recombinations which occur on longer time scales., Here the UVB flux can counteract the recombinations which occur on longer time scales.898 Nevertheless the Sits associated with these galaxies are too small (of order O.LAL. 1) to further increase the transmitted Dux., Nevertheless the SFRs associated with these galaxies are too small (of order $0.1 M_\odot$ $^{-1}$ ) to further increase the transmitted flux.899 Boosting the SER up to values around 50 M. vet (dashed-dotted line). the increase in CP is significant and closely approches the observed value in the innermost region.," Boosting the SFR up to values around 50 $_\odot$ $^{-1}$ (dashed-dotted line), the increase in $\langle F \rangle$ is significant and closely approches the observed value in the innermost region."900 By no means though the flux around 1 Alpe can be matched., By no means though the flux around 1 Mpc can be matched.901 A recent study on the statistics of the Lya forest in the vicinity of foreground. galaxies has revealed an unexpected relative lack of neutral hydrogen in the inner Mpec f+ comoving [rom Ες at 2.~3., A recent study on the statistics of the $\alpha$ forest in the vicinity of foreground galaxies has revealed an unexpected relative lack of neutral hydrogen in the inner Mpc $h^{-1}$ comoving from LBGs at $z \sim 3$.902 High energy. supernova-driven winds able to displace the surrounding σας out to distances greater than 0.57 have been proposed as a viable explanation., High energy supernova-driven winds able to displace the surrounding gas out to distances greater than $0.5 h^{-1}$ have been proposed as a viable explanation.903 However. several numerical. stuclics failed in reproducing the observed clleet of such superwincds via (ADSPLE cosmological simulations.," However, several numerical studies failed in reproducing the observed effect of such superwinds via (M)SPH cosmological simulations."904 According to the simulations. the velocity of the gas closest to galaxies varies around a characterisite value of 40 km s.+ for small objects. and around 150 km + for the biggest ones.," According to the simulations, the velocity of the gas closest to galaxies varies around a characterisitc value of 40 km $^{-1}$ for small objects, and around 150 km $^{-1}$ for the biggest ones."905 These values seem to underestimate the high velocity. gas. at GOO km to measured. by A03 at a scale of a lew kpe from the galaxy center.," These values seem to underestimate the high velocity gas, at 600 km $^{-1}$, measured by A03 at a scale of a few kpc from the galaxy center."906 Nevertheless. the possibility. for the wind to break through the ealactic halo. ancl into the LGAL up to distances greater than Lh comoving Alpe (even at such high velocity) seems quite unphysical if one takes into account the kinetic energy. used to counteract the pressure of the accretion Iow.," Nevertheless, the possibility for the wind to break through the galactic halo, and into the IGM up to distances greater than $1 h^{-1}$ comoving Mpc (even at such high velocity) seems quite unphysical if one takes into account the kinetic energy used to counteract the pressure of the accretion flow."907 Alotivatecl by these reasons. we have studied. the role of local photoionization in determining the Lye absorbers statistics in the vicinity of galaxies. via numerical (ασια(νο transfer simulations.," Motivated by these reasons, we have studied the role of local photoionization in determining the $\alpha$ absorbers statistics in the vicinity of galaxies, via numerical radiative transfer simulations."908 We post-processed an output at 2=3.27 of a multiphase SPL simulation. in which a consistent treatment of supernova feedback. is implemented. running the code to account. for the impact of the UV fux emitted by galaxies on the simulated ICM.," We post-processed an output at $z=3.27$ of a multiphase SPH simulation, in which a consistent treatment of supernova feedback is implemented, running the code to account for the impact of the UV flux emitted by galaxies on the simulated IGM."909 We have derived the synthetic mean Lyo transmitted Εις. as a function of the distance from a typical galaxy. for dilferent. samples of galaxies.," We have derived the synthetic mean $\alpha$ transmitted flux, as a function of the distance from a typical galaxy, for different samples of galaxies."910 We found that on average local photoionization has, We found that on average local photoionization has911returns a Hubble Constant in agreement with the Wey Project.,returns a Hubble Constant in agreement with the Key Project.912 And Scowcrolt (2009) obtained νι = —0.26 in a study of M33., And Scowcroft (2009) obtained $\gamma_{VI}$ = –0.26 in a study of M33.913 Four secondary distance indicators were calibrated by the hey P'oject., Four secondary distance indicators were calibrated by the Key Project.914 The first was the relation (Figure 18)., The first was the Tully-Fisher relation (Figure 18).915 Sakai (2000)obtained Hy = 71 + 3 3+47kms ' F. where the first uncertainty is the random error ancl the second uncertainty is tle systematic error.," Sakai (2000)obtained $_0$ = 71 $\pm$ 3 $\pm$ 7 km $^{-1}$ $^{-1}$, where the first uncertainty is the random error and the second uncertainty is the systematic error."916 Next comes the Fuudameutal plane., Next comes the fundamental plane.917" Ixelsou (2000) used Cepheid clistances to the Leo. Virgo. aud. Fornax clusters to calibrate the fuudamental plane and the D,,.o relation (Figure 20). obtaining Hy = τὸ £5 +49 kins 1 p"," Kelson (2000) used Cepheid distances to the Leo, Virgo, and Fornax clusters to calibrate the fundamental plane and the $_n$ $\sigma$ relation (Figure 20), obtaining $_0$ = 78 $\pm$ 5 $\pm$ 9 km $^{-1}$ $^{-1}$."918 Type Ia supernovae were calibrated by Cibson (2000). using 6 supernova hosts witli measured decline rates. some of them reworked from Saha (1999).," Type Ia supernovae were calibrated by Gibson (2000), using 6 supernova hosts with measured decline rates, some of them reworked from Saha (1999)."919 Figure 21 shows the application of the calibration to supernovae out to 30.000 kins +. vielding Hy = 71 4 2 +47 kin FL ft ," Figure 21 shows the application of the calibration to supernovae out to 30,000 km $^{-1}$, yielding $_0$ = 71 $\pm$ 2 $\pm$ 7 km $^{-1}$ $^{-1}$ ."920Finally. Ferrarese (2000a) calibrated surface brightuess fluctuations iu early type ealaxies. obtaining Hy = 69 £ L£6 kms | f.," Finally, Ferrarese (2000a) calibrated surface brightness fluctuations in early type galaxies, obtaining $_0$ = 69 $\pm$ 4 $\pm$ 6 km $^{-1}$ $^{-1}$ ."921smoothness imposed by the widths of the component eaussians.,smoothness imposed by the widths of the component gaussians.922 It is therefore well-suited. to modeling complex systems like edge-on disc galaxies. where integration along he line of sight. potentially through. multiple components. can make the shape of the LOSVD rather complex.," It is therefore well-suited to modeling complex systems like edge-on disc galaxies, where integration along the line of sight, potentially through multiple components, can make the shape of the LOSVD rather complex."923 In order to increase the signal-to-noise ratio of the resulting kinematic data. the two-dimensional LOSVD as a function of position along the major axis. f(y...2). was assumed to rave the svmmetry of an edge-on axisvmametric dise svstenm. so that Z(eu.A)—F(uuI).," In order to increase the signal-to-noise ratio of the resulting kinematic data, the two-dimensional LOSVD as a function of position along the major axis, $F(v_{\rm los}, R)$, was assumed to have the symmetry of an edge-on axisymmetric disc system, so that $F(v_{\rm924los}, R) \equiv F(-v_{\rm los}, -R)$."925 We therefore determined he kinematic coordinates of the centre of cach galaxy by inding the point in fC...) at which this symmetry. is most closely. obeved.," We therefore determined the kinematic coordinates of the centre of each galaxy by finding the point in $F(v_{\rm los}, R)$ at which this symmetry is most closely obeyed."926 We then averaged the data from the wo sides of each galaxy., We then averaged the data from the two sides of each galaxy.927 The resulting mean estimates for P(r...I) are shown in the top left panels of Figs. 2.. 3..," The resulting mean estimates for $F(v_{\rm los}, R)$ are shown in the top left panels of Figs. \ref{ngc1184plot}, , \ref{ngc1611plot}, \ref{ngc2612plot}, \ref{ngc3986plot},"928 4.. 5.. 6 and , \ref{ngc4179plot} and \ref{ngc5308plot}.929Laving derived these kinematic data. we must now attempt to describe them using a dynamical model.," Having derived these kinematic data, we must now attempt to describe them using a dynamical model."930 The simplest crecible model for the major axis kinematics of a galaxy involved fitting them to a two-integral distribution function. [GEL). which deseribes the phase-space density of stars. where £ is the energy and £ is the angular momentum of the stars about the svnunetry axis. two of the integrals of motion (Binney Tremaine 1987).," The simplest credible model for the major axis kinematics of a galaxy involved fitting them to a two-integral distribution function, $f(E,L)$, which describes the phase-space density of stars, where $E$ is the energy and $L$ is the angular momentum of the stars about the symmetry axis, two of the integrals of motion (Binney Tremaine 1987)."931 Such a moclel would be exact for an infinitelv-thin axisvmmoetric disc. but it is also a credible approximation for a svstem of finite thickness.," Such a model would be exact for an infinitely-thin axisymmetric disc, but it is also a credible approximation for a system of finite thickness."932 Defining he usual polar coordinates. the energy. in the z direction. E.—(Nu)PUB)|407. is approximately an integral of motion for a thin disc.," Defining the usual polar coordinates, the energy in the $z$ direction, $E_z =933\Phi(R,z) - \Phi(R,0) + {1\over2}v_z^2$, is approximately an integral of motion for a thin disc."934 A system whose distribution function akes the form fGF.L)g(£1.) will have dynamics in the z=0 λαο identical to those of an infinitely-thin cise with a distribution function. f/(.L).," A system whose distribution function takes the form $f(E,L)g(E_z)$ will have dynamics in the $z=0$ plane identical to those of an infinitely-thin disc with a distribution function $f(E,L)$."935 Thus. the dynamics in. the ane of the disc can ereclibly be modeled. by treating the major-axis kinematics as if they were those of an infinitely-thin clise.," Thus, the dynamics in the plane of the disc can credibly be modeled by treating the major-axis kinematics as if they were those of an infinitely-thin disc."936 Alathicu Alerrifickl (2000). presented. an algorithm by which the observed density of stars as a function. of projected. radius ancl line-of-sight velocity. CR.ey.) could be iteratively inverted. to find theclistribution function.," Mathieu Merrifield (2000) presented an algorithm by which the observed density of stars as a function of projected radius and line-of-sight velocity, $F(R, v_{\rm los})$ could be iteratively inverted to find thedistribution function,"937(XLE) for LAINBs in the Milky Way in the form of a cut- power law.,(XLF) for LMXBs in the Milky Way in the form of a cut-off power law.938 Since then. Gilfanov. (2004) has produced an average dillerential NLP for LAINBs inaff galaxies.," Since then, Gilfanov (2004) has produced an average differential XLF for LMXBs in galaxies."939 “Phis function is in the form of a power law with two breaks: where Lx=Lx/107 eng and normalizations INquos are related by The bes fit to the overall normalisation is given as In the following we shall use this as the template LAINB XLE for all galaxies.," This function is in the form of a power law with two breaks: where $L_{X,38}=L_X/10^{38}$ erg $^{-1}$ and normalizations $K_{1,2,3}$ are related by The best fit to the overall normalisation is given as In the following we shall use this as the template LMXB XLF for all galaxies."940 Based on Liu et al. (, Based on Liu et al. (9412000. 2001) we estimate the black hole and neutron star fractions amongst the LAINBs to be fgg=0.3. fxs=0.5.,"2000, 2001) we estimate the black hole and neutron star fractions amongst the LMXBs to be $f_{BH} = 0.2$, $f_{NS}942= 0.8$."943 This luminosity function. normalised to the Alilky Way using a total stellar mass of 4.5LOM NL. (Gilfanov 2004) is plotted in Fig 1.," This luminosity function, normalised to the Milky Way using a total stellar mass of $4.5 \times 10^{10}$ $_{\odot}$ (Gilfanov 2004) is plotted in Fig 1."944 For HAINBs. we use the clifferential form of the NLP from Grimm et al. (," For HMXBs, we use the differential form of the XLF from Grimm et al. ("9452003). which is simply a steep power law: where for the Milky Way A~0.7.,"2003), which is simply a steep power law: where for the Milky Way $K \sim 0.7$."946 Dased on Liu οἱ al. (, Based on Liu et al. (9472000. 2001) we estimate the black hole and. neutron star fractions amongst the LAINBs to be fgg=0.05. 0.95.,"2000, 2001) we estimate the black hole and neutron star fractions amongst the HMXBs to be $f_{BH} = 0.05$, $f_{NS} = 0.95$ ."948" ""his function is plotted in Fig 1 alongside the LAINB XLE.", This function is plotted in Fig 1 alongside the LMXB XLF.949 Corbel et al. (, Corbel et al. (9502000: 2003) and Gallo. Fender Pooley (2003) have found an apparently universal correlation between racio and X-ray luminosities for black hole candidate 12116) binaries in the Ἱονπαν and “quiescent? X-ray states (which seem to correspond to accretion rates below about of the Eddington limit).,"2000; 2003) and Gallo, Fender Pooley (2003) have found an apparently universal correlation between radio and X-ray luminosities for black hole candidate (BHC) binaries in the `low/hard' and `quiescent' X-ray states (which seem to correspond to accretion rates below about of the Eddington limit)."951 Phe relation has the form where b~0.7., The relation has the form where $b \sim 0.7$.952 The relation between jet and. racio luminosities may also be considered to have a power-law form Several models of steady. conical jets predict e~LA (e.g. Blancdlord WKonniel 1979: Heinz Sunvacy 2003). which leads to in Edelineton units. or indicating that the fractional jet power increases rapidly as the X-ray luminosity decreases (Fender. Gallo Jonker 2003: Fender. Belloni Gallo 2004).," The relation between jet and radio luminosities may also be considered to have a power-law form Several models of steady, conical jets predict $c \sim 1.4$ (e.g. Blandford Könnigl 1979; Heinz Sunyaev 2003), which leads to in Eddington units, or indicating that the fractional jet power increases rapidly as the X-ray luminosity decreases (Fender, Gallo Jonker 2003; Fender, Belloni Gallo 2004)."953 The value of the normalisation sleetpuc IS uncertain.," The value of the normalisation $A_{\rm steady,BHC}$ is uncertain."954 Fender. Gallo Jonker (2003) esinatcd (in their opinion conservatively) that οωμάBoe=610.," Fender, Gallo Jonker (2003) estimated (in their opinion conservatively) that $A_{\rm955steady, BHC} \geq 6 \times 10^{-3}$."956 While it is certainty [ar from he consensus. larger values of editedBHO may be more realistic.," While it is certainly far from the consensus, larger values of $A_{\rm steady,BHC}$ may be more realistic."957 Alalzac. Moerloni lanan (2004) have suggesed that even at LxI0Lp. re jet may be an order of magnitude more powerful than 10 X-ray emission. which corresponds to Ἑαρμανphoov0.9.," Malzac, Merloni Fabian (2004) have suggested that even at $L_{\rm X} \sim 10^{-3}958L_{\rm Edd}$, the jet may be an order of magnitude more powerful than the X-ray emission, which corresponds to $A_{\rm steady,BHC} \sim9590.3$."960 ernatively. we can set the value of cliopac so that it corresponds to the transition. from. the low/harel to ügh/soft X-ray states (see c.g. AleClintock Itemillard 2005 for a «iscussion). in which case elpHe~0.1.," Alternatively, we can set the value of $A_{\rm steady,BHC}$ so that it corresponds to the transition from the low/hard to high/soft X-ray states (see e.g. McClintock Remillard 2005 for a discussion), in which case $A_{\rm BHC} \sim 0.1$."961 In all the ollowine cliscussions we shall consider 0.006<2ipgce0.3y as covering the full range of likely values., In all the following discussions we shall consider $0.006 \leq A_{\rm BHC} \leq 0.3$ as covering the full range of likely values.962" We furher assume that claseNs=OLehiedeBae. xiwed on the lower ""radio Ioudness"" of NS XRBs (Fender Ixuulkers 2001: Migliari et al."," We further assume that $A_{\rm steady, NS} = 0.1 A_{\rm steady, BHC}$, based on the lower `radio loudness' of NS XRBs (Fender Kuulkers 2001; Migliari et al."963 2003: Fender. Gallo Jonker 2003: Muno et al.," 2003; Fender, Gallo Jonker 2003; Muno et al."964 2004: Alighari Fender in prep)., 2004; Migliari Fender in prep).965 Lt should. be stressed that the LuaxLxUr relation hassof been established for neutron star X-ray binaries. and is only assumed in Fender. Gallo Jonker (2003) by analogy. with the BLUICs.," It should be stressed that the $L_{\rm radio} \propto L_{\rm X}^{0.7}$ relation has been established for neutron star X-ray binaries, and is only assumed in Fender, Gallo Jonker (2003) by analogy with the BHCs."966 1owever. since the BIIC's are so much more racio loud then this is not significant> for the total energy>. budgetὃν calculated here.," However, since the BHCs are so much more radio loud then this is not significant for the total energy budget calculated here."967 A [it to the power in transient optically thin ejection, A fit to the power in transient optically thin ejection968system exacerbates this elfeet.,system exacerbates this effect.969 Llowever because these effects are symmetric the overall global errors remain minimal., However because these effects are symmetric the overall global errors remain minimal.970 This is confirmed. by monitoring the total energv. of the svstem. Figure(4)), This is confirmed by monitoring the total energy of the system. \ref{energy}) )971" traces the fractional variation in total enerey of the system. (£z,dyη."," traces the fractional variation in total energy of the system, $(E_t-E_{t_0})/E_{t_0}$."972 Phe model is à Lowered Evans with 20000 particles. being Tree particles. and timestep df=0.05.," The model is a Lowered Evans with 20000 particles, being Tree particles, and timestep $dt=0.05$."973 Enerey is conserved. in this svstem to within over a period of 50 time units., Energy is conserved in this system to within over a period of 50 time units.974 As ΜΗ all expansion codes angular anc linear momentum are intrinsically not conserved exactly. due to approximations in the force calculation.," As with all expansion codes angular and linear momentum are intrinsically not conserved exactly, due to approximations in the force calculation."975 We might expect i to be more pronounced in the SCETRELE case where the particles in the Tree and SCE codes do not. respond to cach other equally and: oppositely., We might expect it to be more pronounced in the SCFTREE case where the particles in the Tree and SCF codes do not respond to each other equally and oppositely.976 Linear. momentun is known not to be conserved in the pure SCE case. and," Linear momentum is known not to be conserved in the pure SCF case, and"977Ay difference. total sly ditference. (21 %)). and the LOOyam intensity difference at ly: difference (9 %)). respectively.,"$A_{V}$ difference, total $A_{V}$ difference (21 ), and the $100\ \mu m$ intensity difference at $A_{V}$ difference (9 ), respectively."978 On the other hand. the noise of the IRIS data is 0.03 \lJv/sr for A=60jum and 0.06 MJv/sr for A=100µη (Miville Lagache 2005).," On the other hand, the noise of the IRIS data is 0.03 MJy/sr for $\lambda = 60\ \mu m$ and 0.06 MJy/sr for $\lambda = 100\ \mu m$ (Miville Lagache 2005)."979" These values correspond to less than 1.56 for Ay and are negligiblv small compared to the differenο between ;A,: (best) and Ay (SFD98).", These values correspond to less than 1 for $A_{V}$ and are negligibly small compared to the difference between $A_{V}$ (best) and $A_{V}$ (SFD98).980" Figure 10 shows the comparison between 4) (best) and A, (slow).", Figure 10 shows the comparison between $A_{V}$ (best) and $A_{V}$ (slow).981" It can be seen that the difference is large compared to the noise of the DIRBE data points. reflecting the fact that the power of the slow case does not represent the entire correlation in the Cyenus region,"," It can be seen that the difference is large compared to the noise of the DIRBE data points, reflecting the fact that the power of the slow case does not represent the entire correlation in the Cygnus region."982 Figure 11 shows (he comparison between 24\- (best) ancl “dy: (steep)., Figure 11 shows the comparison between $A_{V}$ (best) and $A_{V}$ (steep).983 The difference ls 5% (1 sigma of (chy (best) — Ay (steep))/;dy (best))., The difference is 5 (1 sigma of $A_{V}$ (best) $-$ $A_{V}$ $A_{V}$ (best)).984" The difference between 1\: (steep) and A, (SEDOS) is 20 (1 sigma of (ly (steep) — ely (SED93))/Ay. (steep)).", The difference between $A_{V}$ (steep) and $A_{V}$ (SFD98) is 20 (1 sigma of $A_{V}$ (steep) $-$ $A_{V}$ $A_{V}$ (steep)).985" Thus. the dust temperature difference in the 4, difference is estimated as 18 by Equation (11)."," Thus, the dust temperature difference in the $A_{V}$ difference is estimated as 18 by Equation (11)."986" Figure 12 shows the comparison of ly (steep) and A, (single).", Figure 12 shows the comparison of $A_{V}$ (steep) and $A_{V}$ (single).987 The difference. (24). (steep) — Ay (single))/.Ay. (steep)). scatters by 21 in 1 sigma.," The difference, $A_{V}$ (steep) $-$ $A_{V}$ $A_{V}$ (steep)), scatters by 21 in 1 sigma."988 Ht is consistent with the hvpothesis described in Section 2 that the steep. slow. and best-fit cases are more accurate compared with (he single case.," It is consistent with the hypothesis described in Section 2 that the steep, slow, and best-fit cases are more accurate compared with the single case."989 Dobashi et al. (, Dobashi et al. (990"2005) published an jd, map within |b<40° by the star counting method using the DSS images.",2005) published an $A_{V}$ map within $\mid b \mid \ < 40^\circ$ by the star counting method using the DSS images.991 The spatial resolution is 6'. similar to the present study.," The spatial resolution is $'$ , similar to the present study."992 Figure 13 shows the result of the comparison of both Ay data in the Cygnus region., Figure 13 shows the result of the comparison of both $A_{V}$ data in the Cygnus region.993 The 4A (DSS) data is distributed lower than 24) (best) and saturated at 4 (DSS) > 5 mag., The $A_{V}$ (DSS) data is distributed lower than $A_{V}$ (best) and saturated at $A_{V}$ (DSS) $>$ 5 mag.994 This indicates that the star counting method with DSS has a limit of around Ay (DSS) ~ 5 mag., This indicates that the star counting method with DSS has a limit of around $A_{V}$ (DSS) $\sim$ 5 mag.995 IU is apparent that their method is not useful for Aq > 10 mag because the optical radiation does not reach us because of heavy extinction by the dust., It is apparent that their method is not useful for $A_{V}$ $>$ 10 mag because the optical radiation does not reach us because of heavy extinction by the dust.996 Dobashi (2009) applied their star counting method to the Two-Micron. All Skv Survey (2A\LASS) to reach Ay ~ 30 mag., Dobashi (2009) applied their star counting method to the Two-Micron All Sky Survey (2MASS) to reach $A_{V}$ $\sim$ 30 mag.997 As shown in Section 4.1. Ay bv the best-fit ancl steep cases are more precise than that by SED9S in the Cygnus region.," As shown in Section 4.1, $A_{V}$ by the best-fit and steep cases are more precise than that by SFD98 in the Cygnus region."998 But these methods would not be effective in high Galactic ]atitudes because {ο(140pone) has a lower S/N ratio in those latitudes., But these methods would not be effective in high Galactic latitudes because $I_{C}\ (140\ \mu m)$ has a lower S/N ratio in those latitudes.999 Therefore. if we choose onlv one case out of the best-fit. slow. steep. the main-correlation. and the sub-correlation for the entire sky. the sub-correlation is best suitedbecause of its similarity to the steep case.," Therefore, if we choose only one case out of the best-fit, slow, steep, the main-correlation, and the sub-correlation for the entire sky, the sub-correlation is best suitedbecause of its similarity to the steep case."1000"fractions, from 0 to 0.2 in the primary and satellite galaxies, and a range of orbital parameters.","fractions, from 0 to 0.2 in the primary and satellite galaxies, and a range of orbital parameters."1001" As shown in ?,, minor mergers result in a redistribution of orbital into internal angular momentum, which affects all galaxy components."," As shown in \citet{quDM210a}, minor mergers result in a redistribution of orbital into internal angular momentum, which affects all galaxy components."1002" In particular, old stars, i.e. those already in place before the interaction, always lose angular momentum during the merging process."," In particular, old stars, i.e. those already in place before the interaction, always lose angular momentum during the merging process."1003" The decrease of the specific AM of old stars is accompanied by a redistribution of stellar orbits, as traced by the anisotropy parameter 8, which become increasingly radial."," The decrease of the specific AM of old stars is accompanied by a redistribution of stellar orbits, as traced by the anisotropy parameter $\beta$, which become increasingly radial."1004 In minor mergers with gas in the disk of the primary galaxy we find a similar trend of old stars losing angular momentum and as a result their orbits becoming more radially dominated., In minor mergers with gas in the disk of the primary galaxy we find a similar trend of old stars losing angular momentum and as a result their orbits becoming more radially dominated.1005" However, when a new stellar component forms from gas present in the primary disk during the merger, its AM content is significantly different: the orbits tend to be more tangentially dominated, thus providing a higher rotational support."," However, when a new stellar component forms from gas present in the primary disk during the merger, its AM content is significantly different: the orbits tend to be more tangentially dominated, thus providing a higher rotational support."1006 This different behavior results in a final stellar disk with two different stellar populations with significantly different AM content., This different behavior results in a final stellar disk with two different stellar populations with significantly different AM content.1007" In particular, old stars always show a rotational lag with respect to the young stellar component."," In particular, old stars always show a rotational lag with respect to the young stellar component."1008" If one separates all stars into thin disk stars (at heights |z |x1 kpc from the galaxy midplane) and thick disk stars (at |z |>1 kpc), three different components can be found, with different dynamical properties: (1) young stars in the thin disk, which are rotationally supported and show the highest values of v;, (2) old thin disk stars lagging with respect to the new stars and (3) old thick disk stars lagging with respect to both thin disk components."," If one separates all stars into thin disk stars (at heights $\mid z\mid\le$ 1 kpc from the galaxy midplane) and thick disk stars (at $\mid z\mid >$ 1 kpc), three different components can be found, with different dynamical properties: (1) young stars in the thin disk, which are rotationally supported and show the highest values of $v_t$, (2) old thin disk stars lagging with respect to the new stars and (3) old thick disk stars lagging with respect to both thin disk components."1009" For a minor 1:10 merger, with a satellite accreted on a direct orbit and with an initial primary disk gas fraction of 0.2, the old stars in the thin disk have a rotational lag of about 20 km sl, while the old stars in the thick disk have a velocity about 50 km s! lower than the young stellar component, both lag values being compatible with the estimates for the Milky Way (see ?).."," For a minor 1:10 merger, with a satellite accreted on a direct orbit and with an initial primary disk gas fraction of 0.2, the old stars in the thin disk have a rotational lag of about 20 km $^{-1}$, while the old stars in the thick disk have a velocity about 50 km $^{-1}$ lower than the young stellar component, both lag values being compatible with the estimates for the Milky Way \citep[see ][]{gilmore202}."1010" Multiple mergers can further reduce the tangential velocity of the old stellar components, while leaving that of the new stars mostly unchanged, thus resulting in a further increase in rotational lag with every successive accretion episode."," Multiple mergers can further reduce the tangential velocity of the old stellar components, while leaving that of the new stars mostly unchanged, thus resulting in a further increase in rotational lag with every successive accretion episode."1011" As the two populations, old stars and new stars, have different tangential velocities, in a plot of v, as function of age we expect to find a discontinuity at the time when the merger occurs."," As the two populations, old stars and new stars, have different tangential velocities, in a plot of $v_t$ as function of age we expect to find a discontinuity at the time when the merger occurs."1012" Of course, the newly formed stars will evolve in time, and if no other merger takes place they will be slowly heated by secular effects."," Of course, the newly formed stars will evolve in time, and if no other merger takes place they will be slowly heated by secular effects."1013" However, as we have shown in Fig. 3,,"," However, as we have shown in Fig. \ref{totAMisomer},"1014 secular processes are much less effective in altering stellar kinematics than minor mergers., secular processes are much less effective in altering stellar kinematics than minor mergers.1015" Therefore, in our opinion, a discontinuity in the age-v; plane (or the age-G parameter) between the old and new stellar populations should still be visible, even if secular processes, and asymmetric drift in particular, contribute to the slow heating of the new stellar populations."," Therefore, in our opinion, a discontinuity in the $v_t$ plane (or the $\beta$ parameter) between the old and new stellar populations should still be visible, even if secular processes, and asymmetric drift in particular, contribute to the slow heating of the new stellar populations."1016" Unfortunately, the"," Unfortunately, the"1017In this paper. we report the discovery of a wide VLM binary (hereafter 4445) separated by 130 AU. 373.,"In this paper, we report the discovery of a wide VLM binary (hereafter ) separated by 130 AU, $\farcs$ 3."1018 The brighter primary component of was identified by ?. in the Two Micron All Sky Survey ?) and classified as an M9 dwarf on the ? red optical scheme. indicating a spectrophotometric distance of 33.142.2 pe.," The brighter primary component of was identified by \citet{Reid2008} in the Two Micron All Sky Survey \citep[2MASS;][]{Skrutskie2006} and classified as an M9 dwarf on the \citet{Kirkpatrick1999} red optical scheme, indicating a spectrophotometric distance of $\pm$ 2.2 pc."1019 Neither nor1. activity and age indicators. respectively. were evident in the optical spectrum.," Neither nor, activity and age indicators, respectively, were evident in the optical spectrum."1020 The primaryhas a proper motion of (120414. -25+20) tand a tangential velocity of 19+3 ?).," The primaryhas a proper motion of $\pm$ 14, $\pm$ 20) and a tangential velocity of $\pm$ 3 \citep{Faherty2009}."1021 The system is unresolved in 2MASS. and there have been no reports of a faint companion to this source in either optical survey data or follow-up observations (??)..," The system is unresolved in 2MASS, and there have been no reports of a faint companion to this source in either optical survey data or follow-up observations \citep{Reid2008,1022 Faherty2009}."1023 In our own follow-up observations of4445. we have identified a well-separated. faint L dwarf companion. indicating that this is a wide VLM binary system with a probable BD component.," In our own follow-up observations of, we have identified a well-separated, faint L dwarf companion, indicating that this is a wide VLM binary system with a probable BD component."1024 In Sections ?? and ??.. we describe our imaging and spectroscopic observations. respectively. and discuss the properties of the components of the resolved binary system.," In Sections \ref{Sec: NIRimaging} and \ref{Sec: NIRspectra}, we describe our imaging and spectroscopic observations, respectively, and discuss the properties of the components of the resolved binary system."1025 We discuss the physical association. mass. and age of the binary 4445AB in Section ?? and its implications on VLM formation and evolution scenarios in Section ??..," We discuss the physical association, mass, and age of the binary AB in Section \ref{Sec: analysis} and its implications on VLM formation and evolution scenarios in Section \ref{Sec: discussion}."1026 The conclusions are presented in Section ?2.., The conclusions are presented in Section \ref{Sec: summary}.1027 was imaged with the 3m NASA Infrared Telescope Facility (IRTF) SpeX spectrograph (?) on December 7. 2009 (UT). as part of a program to identify unresolved M/L dwarf plus T dwarf spectral binaries (e.g..?)..," was imaged with the 3m NASA Infrared Telescope Facility (IRTF) SpeX spectrograph \citep{Rayner2003} on December 7, 2009 (UT), as part of a program to identify unresolved M/L dwarf plus T dwarf spectral binaries \citep[e.g.,][]{Burgasser2008a}."1028 Conditions were clear but with poor seeing. 172 at A-band. due in part to the large airmass of the observation (2.34—2.37).," Conditions were clear but with poor seeing, $\farcs$ 2 at $K$ -band, due in part to the large airmass of the observation (2.34–2.37)."1029" These images revealed a faint point source due east of the primary target at à separation of roughly 3"".", These images revealed a faint point source due east of the primary target at a separation of roughly $\arcsec$.1030 Four dithered exposures were obtained of the pair in each of the J. IT. and Jv filters. with individual exposure times of 45s. 30s. and 30s. respectively.," Four dithered exposures were obtained of the pair in each of the $J$, $H$ , and $K$ filters, with individual exposure times of 45s, 30s, and 30s, respectively."1031 The field rotator was aligned at a position angle of0: Le.. north up and east to the left.," The field rotator was aligned at a position angle of $\degr$; i.e., north up and east to the left."1032 Imaging data were reduced in à standard manner using custom IDL routines., Imaging data were reduced in a standard manner using custom IDL routines.1033 Raw images were mirror-flipped about the y-axis to reproduce the sky orientation. and. pair-wise subtracted to remove sky contributions., Raw images were mirror-flipped about the y-axis to reproduce the sky orientation and pair-wise subtracted to remove sky contributions.1034 The difference images were divided by normalized flat field frames. constructed by median-combining the imaging data for each filter after masking out the sources.," The difference images were divided by normalized flat field frames, constructed by median-combining the imaging data for each filter after masking out the sources."1035 Subsections of each image. 10” (83 pixels) on a side and centered on the target source. were extracted from these calibrated frames.," Subsections of each image, $\arcsec$ (83 pixels) on a side and centered on the target source, were extracted from these calibrated frames."1036 A final image for each filter/target pair (Figure 1)) was produced by averaging the registered subframes together. rejecting 5o pixel outhers.," A final image for each filter/target pair (Figure \ref{Fig: image}) ) was produced by averaging the registered subframes together, rejecting $\sigma$ pixel outliers."1037 The two sources of are well resolved along a nearly east-west axis., The two sources of are well resolved along a nearly east-west axis.1038 The brighter western component is hereafter referred toas 4445A and the eastern component as 4445B. Component magnitudes and the angular separation of the pair were determined through point spread function (PSF) fits to the reduced imaging data. following the prescription described in 2..," The brighter western component is hereafter referred to as A and the eastern component as B. Component magnitudes and the angular separation of the pair were determined through point spread function (PSF) fits to the reduced imaging data, following the prescription described in \citet{McElwain2006}."1039 The PSF models were derived from Gaussian fits to the primary component in the individual subimage frames., The PSF models were derived from Gaussian fits to the primary component in the individual subimage frames.1040 For each filter. four distinct PSF models were produced. each of which were fit to the individual images. resulting in a total of 16 independent measures of the relative component magnitudes and 48 independent measures of the separation and orientation. of the pair. in each of the 11 filters.," For each filter, four distinct PSF models were produced, each of which were fit to the individual images, resulting in a total of 16 independent measures of the relative component magnitudes and 48 independent measures of the separation and orientation of the pair, in each of the $JHK$ filters."1041 However. as the secondary was undetected in one of the four J-band images. four measures of the relative J-band flux and separation were discarded before computingmean values and standard deviations.," However, as the secondary was undetected in one of the four $J$ -band images, four measures of the relative $J$ -band flux and separation were discarded before computingmean values and standard deviations."1042 Separation measurements were converted. from pixels to areseconds assuming a plate scale of 01204-07002  (J. Rayner. 2005. private communication) and no distortion.," Separation measurements were converted from pixels to arcseconds assuming a plate scale of $\farcs$ $\pm$ $\farcs$ 002 $^{-1}$ (J. Rayner, 2005, private communication) and no distortion."1043 The position angle (set at 07) was assumed to be accurate to within 0725 (ibid.)., The position angle (set at $\degr$ ) was assumed to be accurate to within $\fdg$ 25 (ibid.).1044 Results are listed in Table 1.., Results are listed in Table \ref{Tab: psf}.1045 The angular separation of the pair was measured to be 37282-07047 at a position angle of 877343:079: r.e.. along an east-west line.," The angular separation of the pair was measured to be $\farcs$ $\pm$ $\farcs$ 047 at a position angle of $\fdg$ $\pm$ $\fdg$ 9; i.e., along an east-west line."1046 The secondary is both considerably fainter and significantly redder than the primary., The secondary is both considerably fainter and significantly redder than the primary.1047 We derived relative magnitudes of A.J= 3.11+0.06 and AA= 2.34+0.04., We derived relative magnitudes of $\Delta{J}=$ $\pm$ 0.06 and $\Delta{K}=$ $\pm$ 0.04.1048 Using the combined-light 2MASS photometry for thesystem?.. this translates into colors of 1.132:0.04 and 1.946 0.08 forthe primary and secondary. respectively.," Using the combined-light 2MASS photometry for the, this translates into colors of $\pm$ 0.04 and $\pm$ 0.08 forthe primary and secondary, respectively."1049 The two components of were observed on separate nights with the prism-dispersed mode of SpeX. the primary on December 7. 2009 (the same night as the imaging observations) and the secondary on December," The two components of were observed on separate nights with the prism-dispersed mode of SpeX, the primary on December 7, 2009 (the same night as the imaging observations) and the secondary on December"1050svstem like Capella Aa.,system like Capella Aa.1051 This also implies that this effect would not be detectable for a slow-rotatine. niain-sequence star like our Sun.," This also implies that this effect would not be detectable for a slow-rotating, main-sequence star like our Sun."1052 Our modeling coufiriis this. showing a total U-I& amplitude of < 0.1 µας for a 1.0 NL... LOR. star witha rotation period of 30.0 days at 10.0 parsecs.," Our modeling confirms this, showing a total U-K amplitude of $\ll$ 0.1 $\mu$ as for a 1.0 $_{\sun}$, 1.0 $_{\sun}$ star with a rotation period of 30.0 days at 10.0 parsecs."1053 These conclusious ou detectabilitv are made with the assmuption that. for bright stars like these. SIM Lite can achieve its 1iicroaresecoud beuchinark.," These conclusions on detectability are made with the assumption that, for bright stars like these, SIM Lite can achieve its microarcsecond benchmark."1054 We show this is possible in narrow angle (NA) mode by cuploving the SIMI Differential Astrometry Performance Estimator (DAPE) (?).., We show this is possible in narrow angle (NA) mode by employing the SIM Differential Astrometry Performance Estimator (DAPE) \citep{Plummer09}.1055 For a target star with magnitude V—5. and a single comparison star with V=L0 located within a deeree of it on the sky. bv iutegratiug 15 seconds on the target. and 30 seconds on the refereuce. for 10 visits at 5 chop cveles cach. a final precision of EL.01 µας is achieved in only 1.01 hours of total nissiou time.," For a target star with magnitude $=$ 5, and a single comparison star with $=$ 10 located within a degree of it on the sky, by integrating 15 seconds on the target, and 30 seconds on the reference, for 10 visits at 5 chop cycles each, a final precision of $\pm$ 1.01 $\mu$ as is achieved in only 1.04 hours of total mission time."1056 For a fainter target with V —10. this precision is ouly reduced to #£1.32 pas iu the same amount of mission time.," For a fainter target with $V$ =10, this precision is only reduced to $\pm$ 1.32 $\mu$ as in the same amount of mission time."1057 Iu utilizine NA mode. one must be careful in choosing the reference star(s). to ensure that they are not stars with a substantial wavelcneth depeudant ceutroid.," In utilizing NA mode, one must be careful in choosing the reference star(s), to ensure that they are not stars with a substantial wavelength dependant centroid."1058 Civeu the ouly coustraiuts on reference stars are that they need to have V z 10 and are within one degree ou the sky. one could easily choose a slow-rotatiug. nmain-sequence star. determined as such via erouud-based observatious. as a waveleugth-audepeudeut astrometric reference star.," Given the only constraints on reference stars are that they need to have V $\gtrsim$ 10 and are within one degree on the sky, one could easily choose a slow-rotating, main-sequence star, determined as such via ground-based observations, as a wavelength-independent astrometric reference star."1059" We also note that wide angle SIM. Lite measurements. witli a precision of —5 µας. may not detect the wavelength dependent photoceter of a system like Capella. but will have no difficulty detecting it in stars like Capella Ab or οσα,"," We also note that wide angle SIM Lite measurements, with a precision of $\sim$ 5 $\mu$ as, may not detect the wavelength dependent photoceter of a system like Capella, but will have no difficulty detecting it in stars like Capella Ab or Vega."1060 The effect of decreasing the eravity darkening exponent is to decrease the total amplitude of the effect in each wavelength. with shorter waveleneths affected more thu longer wavelengths.," The effect of decreasing the gravity darkening exponent is to decrease the total amplitude of the effect in each wavelength, with shorter wavelengths affected more than longer wavelengths."1061 Thus. the choice of gravity darkening exponent is intimately tied to the derived inclination.," Thus, the choice of gravity darkening exponent is intimately tied to the derived inclination."1062 If one were to model observed data with a eravity darkening exponent that was ~LO% different than the true value. they would derive au inclination that would also be ~10% ciffercut from the true inclination.," If one were to model observed data with a gravity darkening exponent that was $\sim$ different than the true value, they would derive an inclination that would also be $\sim$ different from the true inclination."1063 However. the," However, the"1064"The parameters m these models are f,=0.5.0.1. 0.01. fifAL)=U038pCM)/8p(2«10AL...) with δρ as in ((3}). vt=136M... fop=33 Myr. and f,,=0.5. The first του of these parameters appear in combination both iu he expression for the overall streneth of the winds (x fFofav) and the metallicity (x fav).","The parameters in these models are $f_\star = 0.5, 0.1, 0.01$ , $f_w(M)1065= 0.3\delta_B(M)/\delta_B(2 \times 10^8 M_\odot)$ with $\delta_B$ as in \ref{eq:deltaB}) ), $\nu^{-1} = 136 M_\odot$, $t_{\rm OB}= 33$ Myr, and $f_m = 0.5.$ The first three of these parameters appear in combination both in the expression for the overall strength of the winds $\propto f_\star f_w \nu)$ and the metallicity $\propto1066f_\star \nu$ )."1067 Ou the other haud. fop has almost no effect on our results as the relevaut nues for star formation are small compared to structure ormation times scales.," On the other hand, $t_{\rm OB}$ has almost no effect on our results as the relevant times for star formation are small compared to structure formation times scales."1068 Thus we can provide a conservative estimate of the model uncertainties introduced by these xuanieters by simply cousiderius a wide rauge of star ornmation efüciencies. and applying a linear shift iu the final metallicity to cstimate the effect of varvine fie.," Thus we can provide a conservative estimate of the model uncertainties introduced by these parameters by simply considering a wide range of star formation efficiencies, and applying a linear shift in the final metallicity to estimate the effect of varying $f_w$."1069" Finally, while the mass loading parameter f£, las little effect on the overall filliug factor. it is Important for galaxy eedback. and we consider its impact in detail 81123."," Finally, while the mass loading parameter $f_m$ has little effect on the overall filling factor, it is important for galaxy feedback, and we consider its impact in detail 4.3."1070 The inmost obvious. vet perhaps most important feature of 1l is that the filliug factor is always substautially ess than unity. ranging from to at 2=3.," The most obvious, yet perhaps most important feature of 1 is that the filling factor is always substantially less than unity, ranging from to at $z=3$."1071" Note hat these values axe consistent with the 20% curicliueut at 2=| fouud iu nmunuerncal simulations by Thacker. Scannapieco. Davis (2002). using a model simila to our f,=U.l case."," Note that these values are consistent with the $20\%$ enrichment at $z=4$ found in numerical simulations by Thacker, Scannapieco, Davis (2002), using a model similar to our $f_\star = 0.1$ case."1072" The fact that ICAL enrichiueut is inhomogeneous even in the maximal case im which of all barvons in collapsed objects are taken to form stars. however. leads us to au important conclusion: starburst dviven outflows. while au effective source of metals in overdeuse regious (SD). are uot able to eurich the ICAL in its entirety,"," The fact that IGM enrichment is inhomogeneous even in the maximal case in which of all baryons in collapsed objects are taken to form stars, however, leads us to an important conclusion: starburst driven outflows, while an effective source of metals in overdense regions (SB), are not able to enrich the IGM in its entirety."1073 This is true even in the ACDM . model considered in our simulations. in which chwart ealaxics are formed at very lieh redshifts. aud the barvoulc/dark uatter ratio is relatively high. resulting in a large unmuber of stars.," This is true even in the $\Lambda$ CDM model considered in our simulations, in which dwarf galaxies are formed at very high redshifts, and the baryonic/dark matter ratio is relatively high, resulting in a large number of stars."1074 The details of our results depend sensitively ou the uiiminmni nass scale of the galaxies iu our simulation. iowever. Which is set by our minima virial teniperature of 104 K. Tn the central panel of this figure. we plot a series of models iu which no feedback as per ((7)) is miposed. mit instead we allow outflows only from objects above a fixed mass scale.," The details of our results depend sensitively on the minimum mass scale of the galaxies in our simulation, however, which is set by our minimum virial temperature of $10^4$ K. In the central panel of this figure, we plot a series of models in which no feedback as per \ref{eq:strip}) ) is imposed, but instead we allow outflows only from objects above a fixed mass scale."1075 Both the redshift at which outflows ein to become miportanut and their overall filling factor depends closely on this mass., Both the redshift at which outflows begin to become important and their overall filling factor depends closely on this mass.1076" Thus. while iu the ru with f,=0.1. outflowing bubbles fill of the volume at redslift LL12 and reach a final filling factor of 16% exchiding all objects with masses below 1.1«10?A£., shifts these values to z2:S and 6% respectively."," Thus, while in the run with $f_\star = 0.1$, outflowing bubbles fill of the volume at redshift $\lsim 12$ and reach a final filling factor of $16\%$, excluding all objects with masses below $1.1 \times 10^9 M_\odot$ shifts these values to $z \approx 8$ and $6\%$ respectively."1077 Note that this lower resolution is similar to that adopted by Aguirre et ((2001a) and approximately equal to the mass of a sinele dark matter particle in the simulations by Cen aud Ostriker (1990)., Note that this lower resolution is similar to that adopted by Aguirre et (2001a) and approximately equal to the mass of a single dark matter particle in the simulations by Cen and Ostriker (1999).1078 Tn spite of the seusitivitv of metal enrichment to low-uass objects. its overall dependence ou barvouic stripping eedback is weak. as can be seen by comparing the solid ines dn which equation (7)) has been imposed with the dashed lines in which such feedback from outflows is welected.," In spite of the sensitivity of metal enrichment to low-mass objects, its overall dependence on baryonic stripping feedback is weak, as can be seen by comparing the solid lines in which equation \ref{eq:strip}) ) has been imposed with the dashed lines in which such feedback from outflows is neglected."1079 The shape aud fal value of the filline factor are extremely similar between such models for all values of f. )ocomine indistinguishable in niauy cases.," The shape and final value of the filling factor are extremely similar between such models for all values of $f_\star$, becoming indistinguishable in many cases."1080 This is because xuvonie stripping can only occur in a perturbation that is sutiicicutly nearby aud latecollapsing., This is because baryonic stripping can only occur in a perturbation that is sufficiently nearby and late–collapsing.1081 Then the shock velocity. Ry. is large and the overdenuse region occupies a huge solid anele. w. when the outflow reaches it.," Then the shock velocity, $\dot R_s$, is large and the overdense region occupies a large solid angle, $\omega$, when the outflow reaches it."1082 Thus the perturbations succtunbing to barvonic stripping correspond to lateforming galaxies in the most heavily xopulated regions of space. which have little effect ou the overall &lliug factor.," Thus the perturbations succumbing to baryonic stripping correspond to late–forming galaxies in the most heavily populated regions of space, which have little effect on the overall filling factor."1083 The higher bias of suppressed objects can also be secu w comparing the evolution of the filliue factor with the overall mass-averaged ICAL metallicity. plotted in the right xuels of Figure 1..," The higher bias of suppressed objects can also be seen by comparing the evolution of the filling factor with the overall mass-averaged IGM metallicity, plotted in the right panels of Figure \ref{fig:vol}."1084 Iu these panels. the differences between he models with and without suppression are ΙΟ more xonounced.," In these panels, the differences between the models with and without suppression are much more pronounced."1085" The ditfercuce is most apparent iu the f,=4.5 case. im which the wind velocities are the highest. aux lus the suppression of ucighbors is most severe."," The difference is most apparent in the $f_\star=0.5$ case, in which the wind velocities are the highest, and thus the suppression of neighbors is most severe."1086 In this case at 2=3 the overall metallicities differ bv a factor of 1.5 while the difference in volume filliug factor is less than a factor of 1.15., In this case at $z = 3$ the overall metallicities differ by a factor of 1.5 while the difference in volume filling factor is less than a factor of 1.15.1087 Note that the mass-averaged metallicity scales alios huearly with f.. as this parameter controls the ΠΙΟ of stars formed in each galaxy. and hence the ΠΙΟ: of supernovac aud mass of ejected metals.," Note that the mass-averaged metallicity scales almost linearly with $f_\star$, as this parameter controls the number of stars formed in each galaxy, and hence the number of supernovae and mass of ejected metals."1088 We find tha at.=3. ZocmQf. where this relation depends on the assmnued vield (2A/.. per SN. 1/2 ejected). the gas ejected fraction (50%). and the iuinuuni mass scaleim the simulation.," We find that at $z = 3$, $Z \approx 0.1 f_\star$, where this relation depends on the assumed yield $2 M_\odot$ per SN, 1/2 ejected), the gas ejected fraction $50\%$ ), and the minimum mass scalein the simulation."1089" This mass dependence.while seusitive. is more limited than that of the overall filliug factor. as cau be seen by comparing the f,=0.1 model with the series of iiodels with a threshold mass miposed. plotted in the"," This mass dependence,while sensitive, is more limited than that of the overall filling factor, as can be seen by comparing the $f_\star = 0.1$ model with the series of models with a threshold mass imposed, plotted in the"1090stars are aligned (always meaning as seen from the observer with G=0)?,stars are aligned (always meaning as seen from the observer with $G=0$ )?1091 Should we make the recognition by calling the circular image Chwolson ring instead of Einstein ring?, Should we make the recognition by calling the circular image Chwolson ring instead of Einstein ring?1092 We consider a few aspects before casting an intellectually reasonable vote., We consider a few aspects before casting an intellectually reasonable vote.1093same energy band.,same energy band.1094 The brightest region does not correspond to the region with the minimum photon energy. at odd with what we observe in the FilD cloud (Fig. 7..," The brightest region does not correspond to the region with the minimum photon energy, at odd with what we observe in the FilD cloud (Fig. \ref{fig:avgEAB},"1095 right panel)., right panel).1096 In. Paper 1. we singled out spatial regions with homogeneous physical properties and we performed a spatially resolved spectral analysis on them.," In Paper I, we singled out spatial regions with homogeneous physical properties and we performed a spatially resolved spectral analysis on them."1097 In order to compare the observed spectra with those synthesized from the hydrodynamic simulations. we have used the same procedure. by defining physically homogeneous regions.," In order to compare the observed spectra with those synthesized from the hydrodynamic simulations, we have used the same procedure, by defining physically homogeneous regions."1098 The regions are indicated in the upper right panel of Fig. 5..," The regions are indicated in the upper right panel of Fig. \ref{fig:mappeXAB},"1099 for setup Sphl., for setup Sph1.1100 In each of these regions there are limited fluctuations of mean photon energy (S1.7% m region a. €6% in region f. and €2% in region y).," In each of these regions there are limited fluctuations of mean photon energy $\la 1.7\%$ in region $\alpha$, $\la 6\%$ in region $\beta$ , and $\la 2\%$ in region $\gamma$ )."1101 Region « is in the bright northern part of the X-ray knot. where the emission associated to the transmitted shock dominates and the mean photon energy is low. regio P is in the brightest part of the cloud where both transmittec and reflected shocks contribute to the emission. and regior y is located on the South. where we have high values of temperature and of mean photon energy.," Region $\alpha$ is in the bright northern part of the X-ray knot, where the emission associated to the transmitted shock dominates and the mean photon energy is low, region $\beta$ is in the brightest part of the cloud where both transmitted and reflected shocks contribute to the emission, and region $\gamma$ is located on the South, where we have high values of temperature and of mean photon energy."1102 Part of the regions we selected for the analogous spatially. resolved spectral analysis of the data are shown in Fig. I.., Part of the regions we selected for the analogous spatially resolved spectral analysis of the data are shown in Fig. \ref{fig:XMM}.1103 The spectral fittings were performed simultaneously on the synthesized and the observed MOS spectra., The spectral fittings were performed simultaneously on the synthesized and the observed MOS spectra.1104 In agreement with the findings of Paper I. we adopted a MEKAL model of an optically-thin plasma in CIE. with two thermal components. we fixed Nj=1«10°° em. and we left the model Fe abundance free and linked the Ne abundance to it. so as to have (Ne/Nes)/(Fe/Fes)=4.4.," In agreement with the findings of Paper I, we adopted a MEKAL model of an optically-thin plasma in CIE with two thermal components, we fixed $N_{H}=1\times 10^{20}$ $^{-2}$, and we left the model Fe abundance free and linked the Ne abundance to it, so as to have $(Ne/Ne_\odot)/(Fe/Fe_\odot)=4.4$."1105 We added to the model an energy-independent multiplicative factor to take into account the differences in surface brightness between the synthesized and the observed spectra and the different areas of the spectral regions., We added to the model an energy-independent multiplicative factor to take into account the differences in surface brightness between the synthesized and the observed spectra and the different areas of the spectral regions.1106 Our results are summarized in Table 2.., Our results are summarized in Table \ref{tab:spettriAB}.1107 Notice that. since the observed spectra are by themselves well described by this spectral model (as shown in Paper D. the y values in Table 2. can be considered as an indication of the agreement between the observed and the synthesized spectra.," Notice that, since the observed spectra are by themselves well described by this spectral model (as shown in Paper I), the $\chi^{2}$ values in Table \ref{tab:spettriAB} can be considered as an indication of the agreement between the observed and the synthesized spectra."1108 As shown in the table and in the upper panel of Fig. 8..," As shown in the table and in the upper panel of Fig. \ref{fig:spettriA},"1109 there Is a good agreement between the spectrum synthesized in region a and those observed in the FilD regions with low mean photon energy (region 2 and region 4). but if we compare the spectra of region B and region 4 (1. e. the ones with the highest synthesized and observed surface brightness). we have significant differences (see Table 2)).," there is a good agreement between the spectrum synthesized in region $\alpha$ and those observed in the FilD regions with low mean photon energy (region 2 and region 4), but if we compare the spectra of region $\beta$ and region 4 (i. e. the ones with the highest synthesized and observed surface brightness), we have significant differences (see Table \ref{tab:spettriAB}) )."1110 Instead. the spectrum of region 6 has similar spectral features to the one observed in region 7. where the count rate is lower than in region 4. but the mean photon energy is higher.," Instead, the spectrum of region $\beta$ has similar spectral features to the one observed in region 7, where the count rate is lower than in region 4, but the mean photon energy is higher."1111 The spectrum extracted from region y. not shown in the table. is completely different and significantly harder than all the observed spectra (see the lower panel of Fig. 8)).," The spectrum extracted from region $\gamma$, not shown in the table, is completely different and significantly harder than all the observed spectra (see the lower panel of Fig. \ref{fig:spettriA}) ),"1112 even than those extracted form the RegNE cloud. which ts the hardest X-ray emitting region in the EPIC field of view (see Paper I).," even than those extracted form the RegNE cloud, which is the hardest X-ray emitting region in the EPIC field of view (see Paper I)."1113 InSph2.. the post shock temperature of the cloud (1. e. the temperature behind the transmitted shock) is ~10° K (as in setup Sphl). while the cloud density is slightly higher than in setup Sphl (n.€6 em).," In, the post shock temperature of the cloud (i. e. the temperature behind the transmitted shock) is $\sim10^{6}$ K (as in setup Sph1), while the cloud density is slightly higher than in setup Sph1 $n\la 6$ $^{-3}$ )."1114 We synthesized the X-ray count-rate maps and focal plane spectra also for setup Sph2., We synthesized the X-ray count-rate maps and focal plane spectra also for setup Sph2.1115 The emission morphologies in the three energy bands (0.3—0.5 keV. 0.5—| keV. and 0.3—2 keV) are very similar to those of setup Sphl.," The emission morphologies in the three energy bands $0.3-0.5$ keV, $0.5-1$ keV, and $0.3-2$ keV) are very similar to those of setup Sph1."1116 Moreover. the synthesized X-ray emission presents similar spectroscopic features to setup Sphl (see Sect. 3.1)).," Moreover, the synthesized X-ray emission presents similar spectroscopic features to setup Sph1 (see Sect. \ref{A: spherical cloud}) )."1117 However. the global X-ray luminosity of setupSph2 ts too high (more than one order of magnitude) with respect to the observed one. therefore this setup will not be discussed in detail.," However, the global X-ray luminosity of setupSph2 is too high (more than one order of magnitude) with respect to the observed one, therefore this setup will not be discussed in detail."1118for Rez21 ancl. respectively. all 2.ha,"for $\mbox{Re}~z>1$ and, respectively, all $z$."1119"ve Following steps similar to those leading to (161). we have £(2)=—— that may be decomposed in integration over (0.1] and. |1.x). where""her gi(2) :)includes regularization about z2=1 made explicit in (260))(26)) by —L.=>. wherewher (he singulari(v αἱ z=1 in € can be seen (o result [rom the Prime Number Theorem in the form of οντό)—1=o(1) on the basis of the asymptotic behavior of the function 1999)."," Following steps similar to those leading to \ref{EQN_R}) ), we have $\xi(z)=\frac{\pi^\frac{z}{2}}{\Gamma\left(\frac{z}{2}\right)}\int_0^\infty x^{\frac{z}{2}-1}\phi(x) dx,$ that may be decomposed in integration over $(0,1]$ and $[1,\infty)$, where $g_1(z)$ includes regularization about $z=1$ made explicit in \ref{EQN_B0}) ) by $\frac{1}{z-1}$, where the singularity at $z=1$ in $\xi$ can be seen to result from the Prime Number Theorem in the form of $2\sqrt{x}\phi(x)-1=o(1)$ on the basis of the asymptotic behavior of the function \citep{dus99}."1120. The substitution .c=€? gives (25))., The substitution $x=e^{2\lambda}$ gives \ref{EQN_R4}) ).1121 , $\Box$ .1122In a neighborhood of 0<z«1. we may write where wy(2) is analvtic about z>0.," In a neighborhood of $0<z<1$, we may write where $u_1(z)$ is analytic about $z>0$."1123 With z=a4 ib. the second term on the right haud side in the expancecl Evler’s identity (7)) satisfies whereby it is bounded in Re 2=a>5.," With $z=a+ib$ , the second term on the right hand side in the expanded Euler's identity \ref{EQN_B2}) ) satisfies whereby it is bounded in Re $z=a>\frac{1}{2}$."1124 Since the second (term ¢(22) in (7)) is analvGe in Rez=a> 4. it follows that g(a) as defined in Proposition 3.1 is analvtic on a>+.," Since the second term $\zeta(2z)$ in \ref{EQN_B2}) ) is analytic in Re $z=a>\frac{1}{2}$ , it follows that $g(a)$ as defined in Proposition 3.1 is analytic on $a>\frac{1}{2}$."1125 In view ol the analytic and finite behavior of the right hand sidein (7)). the first aud second term on the left hand side in (7)) remain balanced as à approaches 4fromthe right. giving where oli) is analvtic al à= 4.," In view of the analytic and finite behavior of the right hand sidein \ref{EQN_B2}) ), the first and second term on the left hand side in \ref{EQN_B2}) ) remain balanced as $a$ approaches $\frac{1}{2}$fromthe right, giving where $u_2(a)$ is analytic at $a=\frac{1}{2}$ ."1126 As à approaches 4 [rom theright. we have," As $a$ approaches $\frac{1}{2}$ from theright, we have"1127"(2006).. 1£ £i""moc. the cnerev losses. and annihilation rates diller by only a factor of —10 and a substantial fraction of high energy positrons annihilate during the slow-down process.",", if $E_{\rm kin}\gg m_ec^2$, the energy losses and annihilation rates differ by only a factor of $\sim$ 10 and a substantial fraction of high energy positrons annihilate during the slow-down process."1128 These high energy. positrons will form a broad feature at a mean energy cm««|Lys2. while the slowed down positrons will annihilate at. much lower energies and will power a narrow 511 keV line.," These high energy positrons will form a broad feature at a mean energy $\sim1129m_ec^2+E_{\rm kin}/2$, while the slowed down positrons will annihilate at much lower energies and will power a narrow 511 keV line."1130 Thus. a high energy (above 511 keV) component is expected. to be present in the spectrum. with the relative intensity with respect to the narrow 511 keV line depending on the initial positron energy. {μμ and the ionization state of the interstellar medium.," Thus, a high energy (above 511 keV) component is expected to be present in the spectrum, with the relative intensity with respect to the narrow 511 keV line depending on the initial positron energy $E_{\rm kin}$ and the ionization state of the interstellar medium."1131 The latter controls the contribution of the ionization ancl Coulomb losses to the total energy. loss rate., The latter controls the contribution of the ionization and Coulomb losses to the total energy loss rate.1132 Shown in Fig., Shown in Fig.1133 19 is the observed. spectrum. of the julge component and the expected. in-flight annihilation for a neutral (solid) and. ionizecl (dashed) medium., \ref{fig:ia} is the observed spectrum of the Bulge component and the expected in-flight annihilation for a neutral (solid) and ionized (dashed) medium.1134 The clifference in normalization is due to the larger energv losses in the ionized. medium. which increases the fraction of the slowed down positrons at the expense of in-Hight annihilation.," The difference in normalization is due to the larger energy losses in the ionized medium, which increases the fraction of the slowed down positrons at the expense of in-flight annihilation."1135 Phe fraction of slowed down positrons annihilating via positronium formation was set to {ως=0.07., The fraction of slowed down positrons annihilating via positronium formation was set to $f_{\rm ps}=0.97$ .1136 The spectra shown correspond to initial positron energies of 1. 3. 5. 10. 50 ancl 100 MeV. No significant Ilux above 511 keV is observed by SPL in the Bulge component (for the two-component model described in refsecitemplates)). and in Fig.," The spectra shown correspond to initial positron energies of 1, 3, 5, 10, 50 and 100 MeV. No significant flux above 511 keV is observed by SPI in the Bulge component (for the two-component model described in \\ref{sec:templates}) ), and in Fig."1137 LO we show the corresponding 20 upper limits., \ref{fig:ia} we show the corresponding $\sigma$ upper limits.1138 IEvidently. the SPI data above 511 keV do not place tight constraints on the initial energy of positrons.," Evidently, the SPI data above 511 keV do not place tight constraints on the initial energy of positrons."1139" As discussed by Beacom&Yüksel(2006):Sizun.Casse.Schanne(2006).. tighter constraints come from ςΟΛΗTEL data combined. with the SPI measurements of the 511 keV line Hus and. fi. restricting £i, to less than 37.5 MeV. depending on the ionization state of the medium (seeGCregion).."," As discussed by \citet{2006PhRvL..97g1102B,2006PhRvD..74f3514S}, tighter constraints come from COMPTEL data combined with the SPI measurements of the 511 keV line flux and $f_{\rm ps}$, restricting $E_{\rm kin}$ to less than 3–7.5 MeV, depending on the ionization state of the medium \citep[see1140 also][for similar calculations based on earlier measurements of the1141 gamma-ray flux from the GC region]{1981SvAL....7..395A}."1142 The limits on the Bulec Hux above 511. keV clo not stronely constrain the amount of cosmic rays in theDulge. which may produce positrons via a procluction.," The limits on the Bulge flux above 511 keV do not strongly constrain the amount of cosmic rays in theBulge, which may produce positrons via $\pi^+$ production."1143" Indeed. the sae cosmic rays would also produce a comparable amount of z. which would be visible as gamma-ray emission with a peak around LOO MeV. As discussed. by c.g. Abaronian&Atovan (2000). the total gamma-ray [lux from the Inner Galaxy associated. with x"" decav does not exceed ~10photcmτςlay corresponding to a [lux ἳ withing the Dulge area (assuming a solid. angle of the Bulge of —0.1sr."," Indeed, the same cosmic rays would also produce a comparable amount of $\pi^0$, which would be visible as gamma-ray emission with a peak around 100 MeV. As discussed by e.g. \citet[][]{2000A&A...362..937A}, the total gamma-ray flux from the Inner Galaxy associated with $\pi^0$ decay does not exceed $\sim 10^{-4}~{\rm1144 phot~cm^{-2}~s^{-1}~sr^{-1}}$, corresponding to a flux $\lesssim1145 10^{-5}~{\rm phot~cm^{-2}~s^{-1}}$ withing the Bulge area (assuming a solid angle of the Bulge of $\sim 0.1~{\rm sr}$."1146 This is about two orders of magnitude smaller than the observed. Bulge παν 10phots in the annihilation linc., This is about two orders of magnitude smaller than the observed Bulge flux $10^{-3}~{\rm phot~s^{-1}}$ in the annihilation line.1147 The curves shown in Fig., The curves shown in Fig.1148 10. are scaled. by the observed rate of. positron woduction and they approximately match the SPL upper imits (for the initial energv of positrons ~LOO MeV)., \ref{fig:ia} are scaled by the observed rate of positron production and they approximately match the SPI upper limits (for the initial energy of positrons $\sim$ 100 MeV).1149 Therefore. the SPL upper limits on the in-flight annihilation of positrons constrain the mw production rate to be less han ~10sfl.," Therefore, the SPI upper limits on the in-flight annihilation of positrons constrain the $\pi^+$ production rate to be less than $\sim 10^{-3}~{\rm s^{-1}}$."1150 Since the w and x” production rates w the same cosmic ravs are comparable. it is obvious hat observations of LOO MeV. gamma-rays provide much ῃehter constraints on the amount of cosmic ravs in the E2ulge.," Since the $\pi^+$ and $\pi^0$ production rates by the same cosmic rays are comparable, it is obvious that observations of $\sim 100$ MeV gamma-rays provide much tighter constraints on the amount of cosmic rays in the Bulge."1151 We now proceed with spectral fitting of the 511 keV line and ortho-positronium continuum below 511 keV. The simplest possible model is a combination of a Gaussian at 511 keV to describe two-photon annihilation and the three-photon spectrum of Ore&Powell (1949).., We now proceed with spectral fitting of the 511 keV line and ortho-positronium continuum below 511 keV. The simplest possible model is a combination of a Gaussian at 511 keV to describe two-photon annihilation and the three-photon spectrum of \citet{1949PhRv...75.1696O}. .1152 The line normalization. energv and width and the normalization of the ortho- continuum. are free. parameters of the moclel.," The line normalization, energy and width and the normalization of the ortho-positronium continuum are free parameters of the model."1153 The best-fitting values of these parameters are given in Table, The best-fitting values of these parameters are given in Table1154Mapping the entire disk of M31 at mic-inlrarecl wavelengths allows local aid global studies of the galaxy.,Mapping the entire disk of M31 at mid-infrared wavelengths allows local and global studies of the galaxy.1155 Observations with the Infrared Array. Camera (IRAC: on theTelescope simultaneously (trace the dust in the spiral arms at and the oldest stars in the disk and bulge al 3.6 and without the complicating extinction or distance effects that make such studies in the Milky Way difficult., Observations with the Infrared Array Camera \citep[IRAC;][]{irac} on the simultaneously trace the dust in the spiral arms at and the oldest stars in the disk and bulge at 3.6 and without the complicating extinction or distance effects that make such studies in the Milky Way difficult.1156 IRAC observations of M31 are complemented by deep data now available al many other wavelengths., IRAC observations of M31 are complemented by deep data now available at many other wavelengths.1157 They also complement.Spitzer studies of other nearby galaxies., They also complement studies of other nearby galaxies.1158 This paper presents an initial look at the IRAC observations of. M31. focusing on the surface brightness profiles and extended: emission.," This paper presents an initial look at the IRAC observations of M31, focusing on the surface brightness profiles and extended emission."1159 Companion papers discuss longer-wavelength. MIPS observations of M31 (Gordonοἱal.2006).. FRAC and MIPS observations of the M31 satellite galaxy NGC 205 (Marleauetal.2006)... and the implications of the ealaxy’s morphology as seen in non-stellar emission (Blocketal.2006).," Companion papers discuss longer-wavelength MIPS observations of M31 \citep{gordon06}, IRAC and MIPS observations of the M31 satellite galaxy NGC 205 \citep{marleau06}, and the implications of the galaxy's morphology as seen in non-stellar emission \citep{block06}."1160.. A distance to AI3l of 783 kpe (Stanek&Garnavich1998). is assumed throughout.," A distance to M31 of 783 kpc \citep{sg98}1161 is assumed throughout."1162 All magnitudes are on the Vega svstem. using the calibration given bv Reachetal.(2005)..," All magnitudes are on the Vega system, using the calibration given by \citet{reach05}."1163 The IRAC observations of M31 were taken as part ofSpitzer General Observer program 3126 in 2005 January and Fifteen Astronomical Observation Requests (AORs) were used to map a region approximately 37x176 (chosen to match the Spifzer/MIPS observations mace as part of program ID 99). with an extension to the NW to include NGC 205.," The IRAC observations of M31 were taken as part of General Observer program 3126 in 2005 January and Fifteen Astronomical Observation Requests (AORs) were used to map a region approximately $3\fdg7 \times 1\fdg6$ (chosen to match the /MIPS observations made as part of program ID 99), with an extension to the NW to include NGC 205."1164 The central 126x(054 was covered. by three AORs. each having two I2-second [rames per position.," The central $1\fdg6 \times 0\fdg4$ was covered by three AORs, each having two 12-second frames per position."1165 The outer regions were covered bv two AORs each with two clithered 30-second frames per position., The outer regions were covered by two AORs each with two dithered 30-second frames per position.1166 This mapping strategy ensured that observations of each position in the galaxy. were separated by at least 2.5 hours. allowing efficient asteroid rejection in data processing.," This mapping strategy ensured that observations of each position in the galaxy were separated by at least 2.5 hours, allowing efficient asteroid rejection in data processing."1167 The complete dataset consists of 3000 individual images in each of the IRAC channels., The complete dataset consists of 3000 individual images in each of the IRAC channels.1168 Data reduction began with the Basie Calibrated Data (BCD) produced by versions 11 (for the January data) or 12 (he August data) of theSpizer Science Center (SSC) Pipeline., Data reduction began with the Basic Calibrated Data (BCD) produced by versions 11 (for the January data) or 12 (the August data) of the Science Center (SSC) Pipeline.1169" A ""delta. dark’ offset correction was applied to the 12-second frames to correct. for the first-lrame effect. Lollowed by use of the ‘artifact corrector. software developed by S. Carey. which attempts to remove the electronic effects caused by bright stars."," A `delta dark' offset correction was applied to the 12-second frames to correct for the first-frame effect, followed by use of the `artifact corrector' software developed by S. Carey, which attempts to remove the electronic effects caused by bright stars."1170 The remaining, The remaining1171Chicago. Fermilab. the Institute for Advanced Study. the Japan Participation Group. The Johns Hopkins University. Los Alamos National Laboratory. the Max-Planck-Institue for Astronomy (MPIA). the Max-Planck-Institute for Astrrophysics (MPA). New Mexico State University. University of Pitsburgh. Princeton University. the United States Naval Observatory. and the University of Washington.,"Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."1172disk covered by individual structures.,disk covered by individual structures.1173 In this way. we have considered separately the evolution of large active regions and small short-lived ephemeral regions.," In this way, we have considered separately the evolution of large active regions and small short-lived ephemeral regions."1174 Additionally. we have considered the evolution of sunspots (umbrae and penumbrae).," Additionally, we have considered the evolution of sunspots (umbrae and penumbrae)."1175 The coefficients of the neural network are constrained by comparing the output of the model and measurements of the solar irradiance by instruments onboard of SORCE spacecraft., The coefficients of the neural network are constrained by comparing the output of the model and measurements of the solar irradiance by instruments onboard of SORCE spacecraft.1176 The generalization of the network is tested by dividing the data sets on two groups: (1) the training set; and. (2) the validation set.," The generalization of the network is tested by dividing the data sets on two groups: (1) the training set; and, (2) the validation set."1177 We have found that the model error is wavelength dependent., We have found that the model error is wavelength dependent.1178 While the model error for 24-hour forecast in the band from 115 to 180 nm is lower than, While the model error for 24-hour forecast in the band from 115 to 180 nm is lower than1179"spectra from tje three observations we obtaine Lan upoor Li coufideuce) for the fiux of a line at keV dn the backeround spectimm s«—LlJ|κ* photous 97sἘν Ίνα, κτά of the cetected liιο flux.","spectra from the three observations we obtained an upper limit confidence) for the flux of a line at keV in the background spectrum $n < 4.14 \times 10^{-8}$ photons $^{-2}{\rm s^{-1}}$ , i.e. $< 1$ of the detected line flux."1180 The lack of a feaτιre at comparable flux or equivalent widhi n 1ο backeround data also rules out an origin of the kkeV line as a detector feature., The lack of a feature at comparable flux or equivalent width in the background data also rules out an origin of the keV line as a detector feature.1181 We μοι exauined each NIS independenlv. aud ound the KkeV liue to be siguiBcantlv deected iu each NIS unit independenlv. further verifviug he reality of the reported feature (Figure 6)).," We then examined each XIS independently, and found the keV line to be significantly detected in each XIS unit independently, further verifying the reality of the reported feature (Figure \ref{fig:in3xis}) )."1182 Regarding tle more eutative ine at kkeV: the AIS cup does include. a Mn calibration source eiuitine a line at kkeC for the purpose of calibration of the AXIS encYev scale., Regarding the more tentative line at keV: the XIS chip does include a Mn calibration source emitting a line at keV for the purpose of calibration of the XIS energy scale.1183 The calibratiOl sonrco Is ocated at the cu» edge aud the counts from hat source are obviously conservativelv excluled. from ay μ.ο... and bacseround extractio1 cells for scientific alalvsis., The calibration source is located at the chip edge and the counts from that source are obviously conservatively excluded from any source and background extraction cells for scientific analysis.1184" Ilowever awe clo cletec fox weak Mu Ixo emission iu the backerouud προςσπα,"," However, we do detect a weak Mn $\alpha$ emission in the background spectrum."1185 The Mn coutanination Ina a level «10%( |: the measured ne sreneth in the AGN spectruiu., The Mn contamination is at a level $< 10\%$ of the measured line strength in the AGN spectrum.1186" As the source extraction cell is further Yolu the Abu calibratio1 sonrce thaji the backeroui cell and therefore sub,ject to ess Contalatiol. Wwe ονiniate the Mu line COLLaunination frou 1C Caibration source to iIO <LO% for the 20Ef spectrum."," As the source extraction cell is further from the Mn calibration source than the background cell and therefore subject to less contamination, we estimate the Mn line contamination from the calibration source to be $< 10$ for the 2005 spectrum."1187 Tn specral fitting. the exeaest cowcern for the Mu line is the possible contamination ἂν the weak broadered Fe ka CUSSION evideu iu the fit to eieenvector 1 which nav lead to he streneth of the Mu li1ο being over-cstiuated from. simple fits.," In spectral fitting, the greatest concern for the Mn line is the possible contamination by the weak broadened Fe $\alpha$ emission evident in the fit to eigenvector 1 \citep{miller09b} which may lead to the strength of the Mn line being over-estimated from simple fits."1188 Because of the varioρα Issues associated wit la clean lnuüeasureniecut o fiuiv line at keV axd the sesitivity to the conlun form assinned we concentrate onlv on f1ο strong line detection at keV. To confiim the significanceOo of the line we oxforiied Monte Carlo saüulatious usus f1C nethod described iu 7?) and din 7)., Because of the various issues associated with a clean measurement of any line at keV and the sensitivity to the continuum form assumed we concentrate only on the strong line detection at keV. To confirm the significance of the line we performed Monte Carlo simulations using the method described in \citet{porquet04a} and in \citet{markowitz06a}.1189 We took he null hypothesis to be that the προτα is siaplve an absorbed power-law coutimuiun with xuwaneters derived frou fitting the broad-baud data but allowing for the statistical uncertaiutv ou the coutinumun parameters and iucludiug he narrow Fe Ke hue whose presence is well-establisved in this source (2).., We took the null hypothesis to be that the spectrum is simply an absorbed power-law continuum with parameters derived from fitting the broad-band data but allowing for the statistical uncertainty on the continuum parameters and including the narrow Fe $\alpha$ line whose presence is well-established in this source \citep{lobban10a}.1190 We used the cohbunand το create 3000 fake spectra with photon statistics expected from the 2005 exposure. assuming the same iustruments to be operational as for the actual observation.," We used the command to create 3000 fake spectra with photon statistics expected from the 2005 exposure, assuming the same instruments to be operational as for the actual observation."1191 The simulated ¢ata were grouped te» the TWIN cherey-resolution of the iustrumiceus. the same as the observaional data.," The simulated data were grouped to the HWHM energy-resolution of the instruments, the same as the observational data."1192" Following the procedure used to tes the real data for fje presence of a narrow dli1ο, we fitted each fakeo spectruni to obtain the xdues of Ay? obtained roni statistica fluctua10115 in the data."," Following the procedure used to test the real data for the presence of a narrow line, we fitted each fake spectrum to obtain the values of $\Delta \chi^2$ obtained from statistical fluctuations in the data."1193" To map the distribution of AATD across he simulated spectra. each sinimlatec Spectitmo was fitted over the MN3QO keV enerev rango. stepping through using energy bius whose centers were mereased in inercments of LOO eV, The li1 energy was allowed to be free within cach euerev biu teste aud the value of AQ? was recorded at cach point in the spectrum."," To map the distribution of $\Delta \chi^2$ across the simulated spectra, each simulated spectrum was fitted over the 3-10 keV energy range, stepping through using energy bins whose centers were increased in increments of 100 eV. The line energy was allowed to be free within each energy bin tested and the value of $\Delta \chi^2$ was recorded at each point in the spectrum."1194 This nethod makes no assuuptious about the energy at which a lue might be detecteLand over the course of the testing. all energies are tested for he presence of a line.," This method makes no assumptions about the energy at which a line might be detected and over the course of the testing, all energies are tested for the presence of a line."1195 When we fit the simulated da awe are testing wheher we can produce. from statistical fiuctiations. a contribition to X? at he same or greater level as fouix in the actual daa atf any enerev in the range ointerest.," When we fit the simulated data we are testing whether we can produce, from statistical fluctuations, a contribution to $\chi^2$ at the same or greater level as found in the actual data at any energy in the range of interest."1196 The fits to the sinated data vielea distribution of AA? for comparison with the actual data., The fits to the simulated data yielda distribution of $\Delta \chi^2$ for comparison with the actual data.1197" We ound the most extreme statistical fluctuation ο vield a A4? contribution of 19.3 in the LOkkeW baud iu the set of simulated data (νο, πι 3000 simulations no false line appears at any euergy coutributing A?=32 as found in 2005 data).", We found the most extreme statistical fluctuation to yield a $\Delta \chi^2$ contribution of 19.3 in the keV band in the set of simulated data (i.e. in 3000 simulations no false line appears at any energy contributing $\Delta \chi^2=32$ as found in 2005 data).1198" Thus the probability of satisfiusg the null hypothesis is p3:3:«10+ for the line found at keV. As an alcrnative to modeling ταsing individual Caussian lines. we fitted the data uxing a sinele disk line from a narrow annulus. such that the red horn of such a line migi explain the peak at keV “We asstL the system to contain a LnOn-roating black hole. hat the line is Fe Ίνα cuiissio1 from neutra material (at EkkeV). that t1C enüssivitv pattern across the disk cau be deserixd bvor"" where the οἱluissivity iudex q—-2.5 aud thatthe hotsvot exists over n narrow anuuulus of width Ar lr, "," Thus the probability of satisfying the null hypothesis is $p < 3.3 \times 10^{-4}$ for the line found at keV. As an alternative to modeling using individual Gaussian lines, we fitted the data using a single disk line from a narrow annulus, such that the red horn of such a line might explain the peak at keV. We assumed the system to contain a non-rotating black hole, that the line is Fe $\alpha$ emission from neutral material (at keV), that the emissivity pattern across the disk can be described by $r^{-q}$ where the emissivity index q=-2.5 and thatthe hotspot exists over a narrow annulus of width $\Delta r=1r_g$ ."1199We used the same baseline nodel as for testing the Gaussian lines, We used the same baseline model as for testing the Gaussian lines1200as we focus on the stage of the evolution before runaway gas accretion starts.,as we focus on the stage of the evolution before runaway gas accretion starts.

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