ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2model magnitude bin to be composed of both massive. old objects and voung. less massive objects.,"model magnitude bin to be composed of both massive, old objects and young, less massive objects."3 The complication arises at the brightest model bins because the older objects would have masses greater than 0.15... which are not included in the Burrowsetal.(2001) models.," The complication arises at the brightest model bins because the older objects would have masses greater than $M_{\odot}$, which are not included in the \citet{bur} models."4 To account for this. we discard (hose most luminous model bins (compare Figure 3 with Figure 7).," To account for this, we discard those most luminous model bins (compare Figure 3 with Figure 7)."5 The brightest reliable model point is then normalized to the corresponding data point., The brightest reliable model point is then normalized to the corresponding data point.6 The resulting match between the model predictions ancl the data is extremely encouraging. considerimg (hat a rigorous fit to the data was not performed.," The resulting match between the model predictions and the data is extremely encouraging, considering that a rigorous fit to the data was not performed."7 This gives us the confidence to make further predictions about future cluster observations., This gives us the confidence to make further predictions about future cluster observations.8 The voungest of these three clusters. Upper Sco. is at (he voung edge of our models.," The youngest of these three clusters, Upper Sco, is at the young edge of our models."9 The fit to this cluster is «quite good due to the large number of objects detected., The fit to this cluster is quite good due to the large number of objects detected.10 The other two clusters considered. IC: 2391 and the Pleiacles. each have two different age estimates derived using (wo independent techniques (see Table 1).," The other two clusters considered, IC 2391 and the Pleiades, each have two different age estimates derived using two independent techniques (see Table 1)."11 In. both cases. the vounger age estimate (35 and LOO Myr. respectively) is based on matching the turnoff [rom the main-sequence against (he appropriate stellar isochrones (such as Mevnet.Mermilliod.&Maeder (1993))).," In both cases, the younger age estimate (35 and 100 Myr, respectively) is based on matching the turnoff from the main-sequence against the appropriate stellar isochrones (such as \citet{mey93}) )."12 The older age estimate is derived [rom the measurements of the lithium depletion boundary among low-mass cluster members., The older age estimate is derived from the measurements of the lithium depletion boundary among low-mass cluster members.13 Since all of these dwarls are fully convective primordial lithiv is evelecl (through (he core regions and progressively destroyed at a rate which depends on the central temperature., Since all of these dwarfs are fully convective primordial lithium is cycled through the core regions and progressively destroyed at a rate which depends on the central temperature.14 Since the latter parameter depends on mass. lithium is depleted more rapidly in higher mass clwarls.," Since the latter parameter depends on mass, lithium is depleted more rapidly in higher mass dwarfs."15 Thus. as a cluster ages. (he boundary. between clwarls wilh and without detectable lithium absorption moves to lower Iuminosities aud later spectral ivpes.," Thus, as a cluster ages, the boundary between dwarfs with and without detectable lithium absorption moves to lower luminosities and later spectral types."16 Theoretical models can be used to calibrate (hat variation. aud hence estimate cluster ages (Stauffer.Schultz.&Ixirkpatrick.," Theoretical models can be used to calibrate that variation, and hence estimate cluster ages \citep{staf}."171998)... We can compare (he main-sequence filling and lithinm-depletion ages against our estimate based on the predicted morphology of the Iuminosity function., We can compare the main-sequence fitting and lithium-depletion ages against our estimate based on the predicted morphology of the luminosity function.18 In the case of the Pleiades. we cannot distinguish between ages of 100 Myr ancl 120 Avr. since the discriminating features lie 3 to 5 magnitudes fainter than current survey. limit.," In the case of the Pleiades, we cannot distinguish between ages of 100 Myr and 120 Myr, since the discriminating features lie 3 to 5 magnitudes fainter than current survey limit."19 However. the IC 2391 data. are clearly a better match to the 35 Myr luminosity function than the 53 Myr funetion.," However, the IC 2391 data are clearly a better match to the 35 Myr luminosity function than the 53 Myr function."20 The most notable feature in the data is the pronounced peak at Af;14 (Figure 7. middle panel).," The most notable feature in the data is the pronounced peak at $M_I \sim 14$ (Figure 7, middle panel)."21 This corresponds exactly (to the position of peak D (high-mass brown dwarls) in ihe 35 Myr model., This corresponds exactly to the position of peak D (high-mass brown dwarfs) in the 35 Myr model.22 Peak D lies ~2 magnitudes [inter for the 53 Myr age favored by a recent lithium depletion analvsis (DarradovNavascués.Staulfer.&Patten1999)., Peak D lies $\sim2$ magnitudes fainter for the 53 Myr age favored by a recent lithium depletion analysis \citep{bn99}.23. ILowever. the data displaved for IC. 2391 includes both the probable and (he possible members. as defined by DarradovNavascués.Stauffer.&Patten(1999).," However, the data displayed for IC 2391 includes both the probable and the possible members, as defined by \citet{bn99}."24. The majority of the cluster stars contributing to the peak are only possible members. and further observations are required io determine the extent of background contamination.," The majority of the cluster stars contributing to the peak are only possible members, and further observations are required to determine the extent of background contamination."25 To detect features such as peak D. the sample size lor an observed cluster must exceed a," To detect features such as peak D, the sample size for an observed cluster must exceed a"26features of AGN activity.,features of AGN activity.27 A small number of sources display emission lines. but 4 of these are starlorming galaxies based on the NIQAG5S3/La ancl OLLEJASO0 ratios (?)..," A small number of sources display emission lines, but 4 of these are starforming galaxies based on the $\lambda$ $\alpha$ and $\lambda$ $\beta$ ratios \citep{2001ApJ...556..121K}."28 Excluding these. along with all 77the absorption line objects. leaves only 12 (14.8%) sources where an AGN is indicated by the optical spectrum.," Excluding these, along with all the absorption line objects, leaves only 12 $\%$ ) sources where an AGN is indicated by the optical spectrum."29 However. the low spatial resolution of RASS images makes it cüllicult to determine what fraction of X-ray emission originates from the ACN and what fraction may be attributed to the intracluster medium.," However, the low spatial resolution of RASS images makes it difficult to determine what fraction of X-ray emission originates from the AGN and what fraction may be attributed to the intracluster medium."30 This can be compared to both the Lull LASSGAGS catalogue. and the subset of sources which were also observed. as part of the primary survey (2MLASS. selected targets) as shown in ‘Table 6..," This can be compared to both the full RASS--6dFGS catalogue, and the subset of sources which were also observed as part of the primary survey (2MASS selected targets) as shown in Table \ref{clusterspectratab}."31 While less than 15% of the cluster galaxies show optical signatures of an AGN. ~SO of the RASS001ο sources and ~90% of those observed as part of the primary survey (proglD-—1) display emission lines indicative of AGN activity.," While less than $\%$ of the cluster galaxies show optical signatures of an AGN, $\sim 80\%$ of the RASS–6dFGS sources and $\sim 90\%$ of those observed as part of the primary survey (progID=1) display emission lines indicative of AGN activity."32 This confirms that removing the candidate cluster galaxies in the selection process does not eliminate a large number of AGN., This confirms that removing the candidate cluster galaxies in the selection process does not eliminate a large number of AGN.33 By removing these cluster galaxies it also means that we can assume that the majority of the X-ray emission is due to an AGN for our catalogued objects., By removing these cluster galaxies it also means that we can assume that the majority of the X-ray emission is due to an AGN for our catalogued objects.34 77 present a catalogue of X-ray selected AGN [rom theROSA Brieht anc Faint Source. Catalogues with optical spectra in the SDSS Data Release 5.," \citet{2003AJ....126.2209A, 2007AJ....133..313A} present a catalogue of X-ray selected AGN from the Bright and Faint Source Catalogues with optical spectra in the SDSS Data Release 5."35 This catalogue currently covers a magnitude range of 15<g23 and a redshift range of 0.01.<z«4. with median. values of g=15.5 and zo—0.42 respectively.," This catalogue currently covers a magnitude range of $15<g<23$ and a redshift range of $0.01<z<4$, with median values of $g=18.8$ and $z=0.42$ respectively."36 The majority of RASS identifications are quasars and Sevfert 1 galaxies. but there are also a small number of BL-Lac candidates.," The majority of RASS identifications are quasars and Seyfert 1 galaxies, but there are also a small number of BL-Lac candidates."37 Of the 503 objects from this catalogue that are in the southern sky. 126 (25) are also in our RASSGdGS sample.," Of the 503 objects from this catalogue that are in the southern sky, 126 $\%$ ) are also in our RASS–6dFGS sample."38 The RASS Faint Source Catalogue (7). consists of 105.024 sources. making heir combined X-ray. sample more than 6 times larger than he Bright Source Catalogue alone. and accounts for the relatively small fraction of overlap.," The RASS Faint Source Catalogue \citep{2000IAUC.7432....3V} consists of 105,924 sources, making their combined X-ray sample more than 6 times larger than the Bright Source Catalogue alone, and accounts for the relatively small fraction of overlap."39 Figures 12. and 13 clearly show that while the {BSCNVSS is similar to the RASSσος catalogue. the SDSS sample is clistinetively dilferent.," Figures \ref{allzhist} and \ref{allbmaghist} clearly show that while the RBSC–NVSS is similar to the RASS–6dFGS catalogue, the SDSS sample is distinctively different."40 The magnitude in particular is much narrower ancl peaks at a ainter magnitude of g19 while the redshift clistribution »eaks at a sliehthy higher redshift (2~0.2 as opposed. to OL for the other two catalogues)., The magnitude in particular is much narrower and peaks at a fainter magnitude of $g\sim19$ while the redshift distribution peaks at a slightly higher redshift $z\sim0.2$ as opposed to $z\sim0.1$ for the other two catalogues).41 The RASSSDSS catalogue contains fainter. more distant sources resulting rom the fact that the selection. of sources was deeper in oÀh the optical and X-ray regimes.," The RASS–SDSS catalogue contains fainter, more distant sources resulting from the fact that the selection of sources was deeper in both the optical and X-ray regimes."42 Our RASS6dEGS catalogue complements this sample by providing uniform optical spectroscopy for the optically brighter X-ray. sources in the south., Our RASS–6dFGS catalogue complements this sample by providing uniform optical spectroscopy for the optically brighter X-ray sources in the south.43 In this paper we have presented a catalogue of 3405 X-rayselected sources in the Gdb Galaxy Survey., In this paper we have presented a catalogue of 3405 X-rayselected sources in the 6dF Galaxy Survey.44 Sources were selected. from theROSAY All Sky Survey Bright Source Catalogue. hence are count-rate Limited at 0.05 ," Sources were selected from the All Sky Survey Bright Source Catalogue, hence are count-rate limited at $0.05$ $^{-1}$."45Selection criteria of the 6dEPCGS means that these sources are also south of &=OF and outside the Galactic plane (|b|> 10°)., Selection criteria of the 6dFGS means that these sources are also south of $\delta = 0^{\circ}$ and outside the Galactic plane $|b|>10^{\circ}$ ).46 Since the LASS sources were additional targets. not all objects in the catalogue were observed.," Since the RASS sources were additional targets, not all objects in the catalogue were observed."47 However. for completeness. X-ray and available optical information have been included in the catalogue for all targets.," However, for completeness, X-ray and available optical information have been included in the catalogue for all targets."48 A total of 24 (65.34) objects were observed.as part of the survey. with the majority (1715. sources) having reliable redshift measurements.," A total of 2224 $\%$ ) objects were observedas part of the survey, with the majority (1715 sources) having reliable redshift measurements."49 For the optically bright objects. (54 17.5) in the observed. sample. 1333 out. of L478 sources (90.2% ) have reliable redshifts.," For the optically bright objects $b_{\rm J} \leq 17.5$ ) in the observed sample, 1333 out of 1478 sources $\%$ ) have reliable redshifts."50 Inspecting the optical spectra. of the RASSGellCS spectroscopic sample pevealed. that 1171. sources (684) exhibit. broad. emission features. indictative of Type 1AGN. while only 202 (12%) display narrow emission feaures.," Inspecting the optical spectra of the RASS–6dFGS spectroscopic sample revealed that 1171 sources $\%$ ) exhibit broad emission features, indictative of Type 1 AGN, while only 202 $\%$ ) display narrow emission feaures."51 Phe remaining objects are either. absorption-line galaxies (6%) or stars (18%)., The remaining objects are either absorption-line galaxies $\%$ ) or stars $\%$ ).52 The median redshift of the spectroscopic sample is += 0.16. with median redshifts of 2= 0.1706. 2=0.1291 and 2=0.1138 for each of the broad emission. narrow emission and absorption line classes respectively.," The median redshift of the spectroscopic sample is $z=0.16$ , with median redshifts of $z=0.1706$ , $z=0.1291$ and $z=0.1138$ for each of the broad emission, narrow emission and absorption line classes respectively."53the central regions (< 12kpc) have steep negative &eradients. becoming latter at [arger radii.,"the central regions $< 12 \,54\rm kpc$ ) have steep negative gradients, becoming flatter at larger radii."55 High-recdshift (2 2) merger remnants simulated by Wuytsctal.(2010) showed time-dependent negative aas a consequence of redder cores in) redder galaxies (similarly to observations of local galaxies in T|10)., High-redshift $z\sim 2$ ) merger remnants simulated by \cite{Wuyts+10} showed time-dependent negative as a consequence of redder cores in redder galaxies (similarly to observations of local galaxies in T+10).56 eeradients were also predicted. for merger. remnants. by llopkinsetal.(2010).. who stressed that the effective radius derived from the B-band light would increase with ime (bv ~ 10'4)) simply because of stellar. population eradients.," gradients were also predicted for merger remnants by \cite{Hopkins+10}, who stressed that the effective radius derived from the $B$ -band light would increase with time (by $\sim$ ) simply because of stellar population gradients."57 This elfect is important. for conclusions. abou stellar densities. and related: dynamical implications (e.g. or the fundamental plane)., This effect is important for conclusions about stellar densities and related dynamical implications (e.g. for the fundamental plane).58 Similarly vanDokkum(2008) argued that the existence of such gradients can imply tha he observed quiescent. galaxies at large redshift may be of an order or magnitude more dense than previously founc rom observations., Similarly \citet{van Dokkum08} argued that the existence of such gradients can imply that the observed quiescent galaxies at large redshift may be of an order or magnitude more dense than previously found from observations.59 The eeradients present a general complication for estimates of he dark matter distribution in galaxies., The gradients present a general complication for estimates of the dark matter distribution in galaxies.60 Portinari(2010) found that for spiral galaxies. including &eracients would change the fits to circular velocity curves rut not produce strong elfects on the inferred. dark matter density. parameters.," \citet{PS10} found that for spiral galaxies, including gradients would change the fits to circular velocity curves but not produce strong effects on the inferred dark matter density parameters."61 To our knowledge. this issue has not »en addressed in the literature for early-tvpe galaxies. and it is bevond the scope of the present. paper to do so in any detail.," To our knowledge, this issue has not been addressed in the literature for early-type galaxies, and it is beyond the scope of the present paper to do so in any detail."62 In particular. we have studied the stellar population &eracdiecnts only inside 1Aap. while the gradients out to ~ D5 ecould be important for dark matter analwses (e.g. apolitanoetal.2005.2009. 2011)).," In particular, we have studied the stellar population gradients only inside 1, while the gradients out to $\sim$ 5 could be important for dark matter analyses (e.g. \citealt{2005MNRAS.357..691N, Napolitano+09, Napolitano+11}) )."63 Assuming that the &eracients do not change radically outside ~ 1 Roiy.. for the massive EVGs generally. studied with lensing or cvnamics (oy> 150 km 1). we would expect the stellar eeracicntsBS to be mild (see Fig. 3)).," Assuming that the gradients do not change radically outside $\sim$ 1 , for the massive ETGs generally studied with lensing or dynamics $\sig_0 \gsim$ 150 km $^{-1}$ ), we would expect the stellar gradients to be mild (see Fig. \ref{fig:fig3}) ),"64 and to not strongly allect he dark matter estimates., and to not strongly affect the dark matter estimates.65 A final issue is that until now. we have assumed. that he LME is invariant.," A final issue is that until now, we have assumed that the IMF is invariant."66 ΙΓ there are ealaxy-to-ealaxy variations in the IME (see e.g. vanDokkum&Conroynawardhanaetal. 2011)). then the stellar vvalues are changed only by à constant multiplicative value. and the implied eradients remain unchanged.," If there are galaxy-to-galaxy variations in the IMF (see e.g. \citealt{2010Natur.468..940V,vDC11,2011arXiv1104.2379G}) ), then the stellar values are changed only by a constant multiplicative value, and the implied gradients remain unchanged."67 Eowever. EME variations if real would then likely exist.thin. galaxies owing to variations in stellar populations.," However, IMF variations if real would then likely exist galaxies owing to variations in stellar populations."68 Indeed. there have been many suggestions that stars forming at carly times had a non-standard IME (c.g. Larson2005:WKlessenetal.2007:vanDokkum2008:DaveHoldenetal. 2010)).," Indeed, there have been many suggestions that stars forming at early times had a non-standard IMF (e.g. \citealt{2005MNRAS.359..211L,2007MNRAS.374L..29K,van69Dokkum08,2008MNRAS.385..147D,Holden10}) )."70" One study of nearby E/ICis suggested that svstems with older stars within 1 hhave lighter"" (closer to Chabrier than Salpeter) LIMES (Napolitano.Romanowsky&Tortora.|2010).", One study of nearby ETGs suggested that systems with older stars within 1 have “lighter” (closer to Chabrier than Salpeter) IMFs \citep{NRT10}.71.. This would imply that positive ancl negativo age gradients. would decrease and increase the stellar eeracients. respectively.," This would imply that positive and negative age gradients would decrease and increase the stellar gradients, respectively."72 Revisiting Fig. 3..," Revisiting Fig. \ref{fig:fig3},"73 the gradients for the centrallv-xoung objects would decrease and the centralA old. objects would. increase. possibly reducing ancl even eliminating the age-related scatter in the gradient trends.," the gradients for the centrally-young objects would decrease and the centrally old objects would increase, possibly reducing and even eliminating the age-related scatter in the gradient trends."74" Another scenario. would. be if the central regions of massive EPCs have higher-mass"" IMES than standard (vanDokkum&Conroy 2011)). while the outer regions have accreted galaxies with more “normal” IME'Ss (e.g. Ixroupa2001 or Chabrier 2003))."," Another scenario would be if the central regions of massive ETGs have “higher-mass” IMFs than standard \citealt{vDC11}) ), while the outer regions have accreted galaxies with more “normal” IMFs (e.g. \citealt{Kroupa01} or \citealt{Chabrier03}) )."75 This would decrease these galaxies? stellar eeracdients. although it is not clear how this cllect would relate to age and metallicity trends.," This would decrease these galaxies' stellar gradients, although it is not clear how this effect would relate to age and metallicity trends."76 In future analyses. we plan to analyze stellar pprofiles out to a lewRar. and analyze the impact of geracicnts on dark matter profile inferences. which may be compared to ACDAL (e.g. Navarro.Frenk&White 1996)).," In future analyses, we plan to analyze stellar profiles out to a few, and analyze the impact of gradients on dark matter profile inferences, which may be compared to $\Lambda$ CDM (e.g. \citealt{NFW96}) )."77 We will also consider. the (probably. small) adjustments implied for central dark matter inferences (c.g. Tortoraetal. 2000b))., We will also consider the (probably small) adjustments implied for central dark matter inferences (e.g. \citealt{Tortora10lensing}) ).78 Dust effects will further be examined: dust. is not important for metallicity gradients (see T|10) but could allect age and thereby. eeracients. particularly for intermecdiate-mass LCs.," Dust effects will further be examined: dust is not important for metallicity gradients (see T+10) but could affect age and thereby gradients, particularly for intermediate-mass ETGs."79 We thank the referee for his suggestions which helped toimprove the paper., We thank the referee for his suggestions which helped toimprove the paper.80 CE was supported by the Swiss National Science Foundation., CT was supported by the Swiss National Science Foundation.81 AJL was supported by National Science Foundation Grants AST-0808099 ancl AST-0909237., AJR was supported by National Science Foundation Grants AST-0808099 and AST-0909237.82cells.,cells.83 This change in resolution has basically no impact on the disk’s dynamics., This change in resolution has basically no impact on the disk's dynamics.84" As before, the disk eccentricity settles to a time averaged value of about 0.04."," As before, the disk eccentricity settles to a time averaged value of about $ 0.04 $."85 In all three simulations the disk periastron does not reach a state with real precession and the only change caused by the change of resolution is a small shift in time., In all three simulations the disk periastron does not reach a state with real precession and the only change caused by the change of resolution is a small shift in time.86 Consequently our standard resolution of 256x574 cells seems to be sufficient to resolve the disk dynamics properly., Consequently our standard resolution of $ 256 \times 574 $ cells seems to be sufficient to resolve the disk dynamics properly.87" To study the influence of the physical conditions on the evolution of the disk, we investigated in particular the impact of the disk mass Μαι, the viscosity v, which is determined by a, the opacity κ. and the binary's eccentricity epis."," To study the influence of the physical conditions on the evolution of the disk, we investigated in particular the impact of the disk mass $ M_\mathrm{disk} $, the viscosity $ \nu $ , which is determined by $ \alpha $, the opacity $ \kappa $, and the binary's eccentricity $ e_\mathrm{bin} $."88 The first parameter we investigated is the disk mass., The first parameter we investigated is the disk mass.89 In contrast to isothermal simulations our radiative simulations depend on the disk mass as the opacity depends on gas density., In contrast to isothermal simulations our radiative simulations depend on the disk mass as the opacity depends on gas density.90" To analyze the influence of the disk mass on the disk's evolution, we ran four simulations with disk masses Μάικ of 0.005Mo, 0.01Mo, 0.02Mc, and 0.04Mc; while keeping all other parameters unchanged."," To analyze the influence of the disk mass on the disk's evolution, we ran four simulations with disk masses $ M_\mathrm{disk} $ of $ 0.005\,M_{\sun}$, $0.01 \,M_{\sun}$, $0.02\,M_{\sun}$ and $ 0.04 \,M_{\sun} $ while keeping all other parameters unchanged."91" The surface density and temperature profiles after 100 binary orbits for different disk masses is displayed in Fig. 7,,"," The surface density and temperature profiles after $100$ binary orbits for different disk masses is displayed in Fig. \ref{fig:gc-mass-sig-temp},"92 where the solid (red) line refers to the standard model displayed in Fig. 3.., where the solid (red) line refers to the standard model displayed in Fig. \ref{fig:gc-times-sig-temp-ecc}.93 The profiles can be divided into two regimes that are separated by the 1000 KK temperature line caused by the opacity (see Section 3.2))., The profiles can be divided into two regimes that are separated by the $ 1000$ K temperature line caused by the opacity (see Section \ref{sec:structureanddynamics}) ).94" In each regime the surface density and temperature profiles follow a simple power-law that depends only weakly on the disk mass, and is given in the caption of Fig. 7.."," In each regime the surface density and temperature profiles follow a simple power-law that depends only weakly on the disk mass, and is given in the caption of Fig. \ref{fig:gc-mass-sig-temp}."95 Fig., Fig.96 8 shows the time evolution of the disk eccentricity and periastron for different disk masses., \ref{fig:gc-mass-ecc-peri} shows the time evolution of the disk eccentricity and periastron for different disk masses.97" The higher the disk mass, the lower the oscillations of the the disk eccentricity, but the time average of the disk eccentricity is in the range of 0.04 — 0.05 for all disk masses."," The higher the disk mass, the lower the oscillations of the the disk eccentricity, but the time average of the disk eccentricity is in the range of $ 0.04 $ – $ 0.05 $ for all disk masses."98 The disk periastron displays no real precession and with increasing disk mass it settles at about 0., The disk periastron displays no real precession and with increasing disk mass it settles at about $ 0 $ .99 This is in contrast to the isothermal simulations presented in Section 3.1 where we obtained a disk eccentricity egisk of 0.2 for an aspect ratio H/r of 0.05 and a real precession of the disk periastron., This is in contrast to the isothermal simulations presented in Section \ref{sec:isothermal} where we obtained a disk eccentricity $ e_\mathrm{disk} $ of $ 0.2 $ for an aspect ratio $ H/r $ of $ 0.05 $ and a real precession of the disk periastron.100" There is a trend, however, that the eccentricity becomes higher for cooler disks with lower H/r."," There is a trend, however, that the eccentricity becomes higher for cooler disks with lower $H/r$ ."101" But one has to be careful here, because the aspect ratio is constant in radius and time for the isothermal simulations, but not for our radiative simulations."," But one has to be careful here, because the aspect ratio is constant in radius and time for the isothermal simulations, but not for our radiative simulations."102 Fig., Fig.103 9 shows the aspect ratio of the disk after 100 binary orbits., \ref{fig:gc-hr} shows the aspect ratio of the disk after $ 100 $ binary orbits.104 All models had an initial value of (H/r)ii;4;=0.05 but end up with different values of H/r depending on the disk mass and the phase in the binary orbit., All models had an initial value of $ \left( H/r \right)_\mathrm{initial} = 0.05 $ but end up with different values of $ H/r $ depending on the disk mass and the phase in the binary orbit.105 We note that in all models the mass of the disk reduces with time owing tothe mass loss across the outer boundary., We note that in all models the mass of the disk reduces with time owing tothe mass loss across the outer boundary.106" In particular, we find that after 100 binary orbits the disks have lost about 27% of their initial mass in the Μαι=0.04Mo model, 2346 in the Maia;=0.04Mo model, 17% in the Maik=0.04Mo and 6% in the Maia;=0.005Me model."," In particular, we find that after $ 100 $ binary orbits the disks have lost about $ 27\,\% $ of their initial mass in the $ M_\mathrm{disk} = 0.04\,M_{\sun} $ model, $ 23\,\% $ in the $ M_\mathrm{disk} = 0.04\,M_{\sun} $ model, $ 17\,\%$ in the $ M_\mathrm{disk} = 0.04\,M_{\sun} $ and $ 6\,\% $ in the $ M_\mathrm{disk} = 0.005\,M_{\sun} $ model."107" Hence, the values quoted in the text and figures always refer to the initial disk masses."," Hence, the values quoted in the text and figures always refer to the initial disk masses."108" The results for eg, in Fig.", The results for $e_{\rm disk}$ in Fig.109 8 show a marginal increase of the disk eccentricity for smaller disk masses., \ref{fig:gc-mass-ecc-peri} show a marginal increase of the disk eccentricity for smaller disk masses.110" To test if this trend continues, we performed additional simulations for even smaller initial disk masses and found indeed an increased oscillatory behavior of the eccentricity, which settles eventually after about 600 binary orbits to a low eccentric state very similar to the Maia,=0.025 model, however."," To test if this trend continues, we performed additional simulations for even smaller initial disk masses and found indeed an increased oscillatory behavior of the eccentricity, which settles eventually after about $600$ binary orbits to a low eccentric state very similar to the $M_{\rm disk}= 0.025$ model, however."111" Hence, there does not seem to exist an obvious trend of egg, with disk mass."," Hence, there does not seem to exist an obvious trend of $e_{\rm disk}$ with disk mass."112 In this section we investigate the influence of the viscosity v., In this section we investigate the influence of the viscosity $ \nu $.113" To study the dependence on v we varied the a parameter, which determines the viscosity, v=ac,Hacid (?),, from our standard model (a= 0.01) and Κερί all other parameters unchanged."," To study the dependence on $ \nu $ we varied the $ \alpha $ parameter, which determines the viscosity, $ \nu = \alpha c_\mathrm{s} H = \alpha c_\mathrm{s}^2 \Omega_K^{-1} $ \citep{1973A&A....24..337S}, , from our standard model $\alpha = 0.01 $ ) and kept all other parameters unchanged."114 We varied α from 0.005 to 0.04., We varied $ \alpha $ from $ 0.005 $ to $ 0.04 $.115 Fig., Fig.116 10 shows the surface density and temperature profiles after 100 binary orbits., \ref{fig:gc-visc-sig-temp} shows the surface density and temperature profiles after $ 100 $ binary orbits.117" All models started with identical disk mass, and evidently, higher o models lose the disk masses more rapidly."," All models started with identical disk mass, and evidently, higher $ \alpha $ models lose the disk masses more rapidly."118" For example, atthe displayed time at 100 binary orbits the a=0.04 model has lost about 71% ofits initial mass, while the standard model (α= 0.01) lost only 19%, see also Fig."," For example, atthe displayed time at $100$ binary orbits the $ \alpha=0.04$ model has lost about $71\,\%$ ofits initial mass, while the standard model $\alpha=0.01$ ) lost only $19\,\%$ , see also Fig."119 12 below., \ref{fig:gc-visc-mass} below.120" The shape of the surface density and temperature profile seems tobe independent of the viscosity in the disk, as expected."," The shape of the surface density and temperature profile seems tobe independent of the viscosity in the disk, as expected."121" Again,the profiles can be divided into"," Again,the profiles can be divided into"122We define here the model comparison techuique using the fields approach. which has uutil now not been used in the literature.,"We define here the model comparison technique using the fields approach, which has until now not been used in the literature."123 Using a generalised likelihood ratio approach. the quantity to measure is: Theoretically A couvereeuces asviuptotically to a 4? variable with a certain nuuber of defs. As we oulv iive one observation we cannot assume (asviptotic) convergence.," Using a generalised likelihood ratio approach, the quantity to measure is: Theoretically $K_6$ convergences asymptotically to a $\chi^2$ variable with a certain number of 's. As we only have one observation we cannot assume (asymptotic) convergence."124 We can however use a Mone Carlo approach in order to estimate the p-value of the est under the Ly wpothesis., We can however use a Monte Carlo approach in order to estimate the $p$ -value of the test under the $H_0$ hypothesis.125 The p-value is defined as the probahiitv that under Lf) he test value can be over a given As. be. P(f>Ke)= pice. where p is the probability cistribution of the nest under the 11 ivpothesis.," The $p$ -value is defined as the probability that under $H_0$ the test value can be over a given $K_6$, i.e. $P( t >126K_6) = \int_{K_6}^{\infty} p(x) dx$ , where $p$ is the probability distribution of the test under the $H_0$ hypothesis."127 By simulating primordial CXMB for a fducial cosmology. these values cau be easily computed.," By simulating primordial CMB for a fiducial cosmology, these values can be easily computed."128 Then the p-value gives us a confidence ou rejecting the fy hypothesis., Then the $p$ -value gives us a confidence on rejecting the $H_0$ hypothesis.129" Notice hat the same procedure for the ντατο estimation can be applied on A, (Equation. 22)). even ou A4 (Equation 15)) and A; (Equation 17)) for the power spectra micthods."," Notice that the same procedure for the$p$ -value estimation can be applied on $K_4$ (Equation \ref{eq:16}) ), even on $K_1$ (Equation \ref{eq:15}) ) and $K_3$ (Equation \ref{eq:5}) ) for the power spectra methods."130 We will further refer to thi119 p-value estimation as the Alonte-Carlo estimation. as we theoretically know the distribution under the null hypothesis. ic. the primordial CMD is supposed to coue from a Gaussian random process.," We will further refer to this $p$ -value estimation as the Monte-Carlo estimation, as we theoretically know the distribution under the null hypothesis, i.e. the primordial CMB is supposed to come from a Gaussian random process."131 Tn order to validate the Saclay method. we estimate the detection level expected using WALAP 7 data for the CAIB and 2\TASS and Euclid data for the ealaxy data (see section 6 for a description of WALAP aud 2\TASS data sets).," In order to validate the Saclay method, we estimate the detection level expected using WMAP 7 data for the CMB and 2MASS and Euclid data for the galaxy data (see section \ref{sec:data} for a description of WMAP and 2MASS data sets)."132 We quantity the effect of the impainting process ou CMD maps with aud without an ISW signal., We quantify the effect of the inpainting process on CMB maps with and without an ISW signal.133 We do this bv simulating 2MASS-like aud Enclid-like Cussiau and lognormal galaxy distributions aud WALAP7-like Cassia CAD maps (using cosmological parameers frou Table 63) bohi with aud wihout an ISW signal., We do this by simulating 2MASS-like and Euclid-like Gaussian and lognormal galaxy distributions and WMAP7-like Gaussian CMB maps (using cosmological parameters from Table \ref{tab:wmap7}) ) both with and without an ISW signal.134 We then iyply our method to attempt a deection of he ISW sjenal., We then apply our method to attempt a detection of the ISW signal.135 We «o this both on full-sky aps as well as on masked data where we have reconstructed data beline the mask using the sparse inpaiutiug techique (the masks we use are as described iu ?7 aud ??))., We do this both on full-sky maps as well as on masked data where we have reconstructed data behind the mask using the sparse inpainting technique (the masks we use are as described in \ref{sec:data:wmap} and \ref{sec:data:2mass}) ).136 For cach simulation. we run the 3-step Saclay inethod.," For each simulation, we run the 3-step Saclay method."137 Except for the cross-correlation method where we use 100 iterations for the 2ALASS-like auc 1000 for Euclid-like simulations for the j-value estimation aud 201 for the variance estimation (nestec bootstrap). every other Monte-Carlo process was performed using 10 000 iterations All tests were performed iide the spherical harmonics doniun with (c[2.100] for 221ASS aud [2.350] for a Enclid-like survey.," Except for the cross-correlation method where we use 100 iterations for the 2MASS-like and 1000 for Euclid-like simulations for the $p$ -value estimation and 201 for the variance estimation (nested bootstrap), every other Monte-Carlo process was performed using 10 000 iterations All tests were performed inside the spherical harmonics domain with ${\ell \in138 [2,100]}$ for 2MASS and $\ell=[2,350]$ for a Euclid-like survey."139" For the Eyclid-like. 3rvov. wo consider a ealaxy distribution as defined im?).. with mean recshift ;,,=0.5 and slopes 4Dj=1.5."," For the Euclid-like survey, we consider a galaxy distribution as defined in, with mean redshift $z_m=0.8$ and slopes $\alpha = 2, \beta=1.5$."140 We recoistruct the ISW effect created w the projected galaxy istribution of the Euclid surv. bv considering oulv one larec redshift biu.," We reconstruct the ISW effect created by the projected galaxy distribution of the Euclid survey, by considering only one large redshift bin."141 Iu the future. it could be possible to refine such a reconstructec map bv considering tomeeraplic bius. or using inform.ation from the spectroscopic survey.," In the future, it could be possible to refine such a reconstructed map by considering tomographic bins, or using information from the spectroscopic survey."142 As sky coverage maps are no vet available for Euclid. we consider the sanie mask as or 2ATASS aud inpaint regions with missing data following Section 6.3..," As sky coverage maps are not yet available for Euclid, we consider the same mask as for 2MASS and inpaint regions with missing data following Section \ref{sec:data:inpainting}."143 We clioose do to this. rather than simply assuiie a value for the fraction of sky covered (fil). so as to consider more realistic problems relating to the shape of the mask. and to test our inpainting method.," We choose do to this, rather than simply assume a value for the fraction of sky covered $f_{\rm sky}$ ), so as to consider more realistic problems relating to the shape of the mask, and to test our inpainting method."144 The expected detection leves (in units of 7) are reported in Table 3 (2MASS) aud Tade E (Euclid)., The expected detection levels (in units of $\sigma$ ) are reported in Table \ref{tab:simu2mass} (2MASS) and Table \ref{tab:simuEuclid} (Euclid).145 Methods 13 correspondto the 3-step method described in Figure 2.. where (b) aud (ALC) denote bootstrap aud Monte Carlo evaluations of the variance axd the p-value of the test.," Methods $1-3$ correspondto the 3-step method described in Figure \ref{fig:saclaymethod}, where $(b)$ and $(MC)$ denote bootstrap and Monte Carlo evaluations of the variance and the $p$ -value of the test."146 The p-values are converted as a σ value using the following formula: WwWrere pis the p-value. s the correspouding e-score aud erf 14ie inverse error function.," The $p$ -values are converted as a $\sigma$ value using the following formula: where $p$ is the $p$ -value, $s$ the corresponding $\sigma$ -score and $\mathrm{erf}^{-1}$ the inverse error function."147 Th| 2MASS simulations (Table 3)) show that we expect he same level of significance for an ISW detection. whether the mass tracer follows a Cussiau or a lognormal cüstrilition.," The 2MASS simulations (Table \ref{tab:simu2mass}) ) show that we expect the same level of significance for an ISW detection, whether the mass tracer follows a Gaussian or a lognormal distribution."148 In either case the significance is low. around loX lo.," In either case the significance is low, around $~1 \sigma \pm 1\sigma$."149"Tus neans that for 2MASS-like survey, we have a signal to woise ratio (S/N) around 1o."," This means that for 2MASS-like survey, we have a signal to noise ratio (S/N) around $1\sigma$."150 We see uo major difference )otween the expected detection levels of AP and M3., We see no major difference between the expected detection levels of M2 and M3.151 The oulv difference is for M1 (but there is still agreemen with M2 and M3 within 16 error bars) - this may be due ο the fact that the boostrap technique is nore efficient for Gaussian assiunuiptions., The only difference is for M1 (but there is still agreement with M2 and M3 within $1\sigma$ error bars) - this may be due to the fact that the bootstrap technique is more efficient for Gaussian assumptions.152 We also apaly our iiethod to CMD suilations with uo ISW sigual preseit (2 left columns of Table 33). aud fiud a lower detection significance than when an ISW sienal is present.," We also apply our method to CMB simulations with no ISW signal present (2 left columns of Table \ref{tab:simu2mass}) ), and find a lower detection significance than when an ISW signal is present."153 This is true even when mpaiuing ds used to recover missing data. showing that the mpaiutiug method does not introdiice spurious correlations.," This is true even when inpainting is used to recover missing data, showing that the inpainting method does not introduce spurious correlations."154 Tn any case. all mehods suggest it is cμοιIt to detect the ISW signal with Πο] significance usi18o he 2MASS data as a local tracer of the matter distribution.," In any case, all methods suggest it is difficult to detect the ISW signal with high significance using the 2MASS data as a local tracer of the matter distribution."155" Iu Table 1.. we show that an Euclid-like πιrvev. which is optimally designed for au ISW detection ?7).. permüts a nmli higher detection than with à 2MAÀSS SUPVCN,"," In Table \ref{tab:simuEuclid}, , we show that an Euclid-like survey, which is optimally designed for an ISW detection , permits a much higher detection than with a 2MASS survey."156 As with 2\TASS siuulations. we notice that AR and M53 return similar detection levels. which are," As with 2MASS simulations, we notice that M2 and M3 return similar detection levels, which are"157The measurements of the secondary. component. are not svmametrie to those of the primary. and it is impossible to derive a SB2 orbit.,"The measurements of the secondary component are not symmetric to those of the primary, and it is impossible to derive a SB2 orbit."158 Finally. a part of a long period orbit is visible in the large residuals of the SBI orbit.," Finally, a part of a long period orbit is visible in the large residuals of the SB1 orbit."159 Therefore. he svstem could be quadruple. although the CCL of a sole spectrum exhibits one dip only.," Therefore, the system could be quadruple, although the CCF of a sole spectrum exhibits one dip only."160 = DD |17 493p., = BD +17 493p.161 Orbit calculated. ciscarding the + xended: measurements., Orbit calculated discarding the 4 blended measurements.162 When they are taken into account. he blend coellicient is Cy=0.624+0.042. but the orbital elements are not improved.," When they are taken into account, the blend coefficient is $C_0=0.624 \pm 0.042$, but the orbital elements are not improved."163 = BD |20 511., = BD +20 511.164 The measurements confirm he variability of the RY., The measurements confirm the variability of the RV.165" A possible orbit was found with the ollowing ""—(3+125) davs. 75=2446400 JD. eNM=0.5+m d=m(26.70.5) km t and (4854+1.8) kms ή>. it is very uncertain. due to the large errors of the measurements."," A possible orbit was found with the following elements: $P=(3070 \pm 125)$ days, $T_0=2446400 \pm 300$ JD, $e=0.5 166\pm 0.3$, $V_0=( 26.7 \pm 0.5)$ km $^{-1}$ and $K_1 =(4.8 \pm 1.8)$ km $^{-1}$, but it is very uncertain, due to the large errors of the measurements."167 = ID 23158., = HD 23158.168 A FS V type starrather cillieult Oo measure withCORAVEL. with verv lage uncertainties (|—2.7 km ," A F5 V type starrather difficult to measure with, with very large uncertainties $I=2.7$ km $^{-1}$ )."169The possible variability is due to the measurement of JD 2448245: when it is discarded. ο)=12.2 ," The possible variability is due to the measurement of JD 2448245; when it is discarded, $P(\chi^2)=12.2$ ."170Therefore. the variability of the star is not certain.," Therefore, the variability of the star is not certain."171 If it is constant. the RV of the star is =(1.9541Ls1) kms+ = WD," If it is constant, the RV of the star is $\bar{V}=(-1.95 \pm 1.81)$ km $^{-1}$."172 around27635., = HD 27635.173 The S measurements of the secondary seem fixed -28 kms ancl we prefer to discard then.," The 8 measurements of the secondary seem fixed around -28 km $^{-1}$, and we prefer to discard them."174 Otherwise. a SB2 a is obtained with A»=s.3 kms!.," Otherwise, a SB2 orbit is obtained with $K_2=8.3$ km $^{-1}$."175 = LD 23597, = HD 285970.176 A first orbit was published by Gunn (1981).., A first orbit was published by \citet{GrifGun}.177 = IID 33155., = HD 33185.178 X bright SB2 6.67 mag) with a semi- axis expected around 58 mas. which should be easily separatec.," A bright SB2 (6.67 mag) with a semi-major axis expected around 58 mas, which should be easily separated."179 = LD 59450., = HD 59450.180 The star belongs to a triple CPM system and is also 240A. = LID solol = ADS 728SAD., The star belongs to a triple CPM system and is also 2:49A. = HD 80101 = ADS 7288AB.181 A visual binary svstem with separation 0.3 aresec., A visual binary system with separation 0.3 arcsec.182 Phe A component is the SBI with orbital clements in Table 4.., The A component is the SB1 with orbital elements in Table \ref{tab:orb1}.183 Ehe dip of the D component is visible on 9 CCL. with the fixed velocity Vg=(52.08=0.33) kn s ," The dip of the B component is visible on 9 CCF, with the fixed velocity $V_B=(52.08 \pm 0.33)$ km s $^{-1}$ ."184Twenty-nine blended. BV. refer to components A and D. = LID S1997., Twenty-nine blended RV refer to components A and B. = HD 81997.185 Revision of the orbit of Duquennov (1901)., Revision of the orbit of \citet{DM91}.186. = LD 89745, = HD 89745.187 A correction of 0.489 km | was added to the 7 RY measurements derived. [rom for cach component. in order to get the best fit.," A correction of 0.489 km $^{-1}$ was added to the 7 RV measurements derived from for each component, in order to get the best fit."188 = LUD 92787., = HD 92787.189 A FS star with laree RY errors., A F5 star with large RV errors.190 A possible secondary component was detected on one correlation dip. and it is possible that. the. other measurements contain blended observations.," A possible secondary component was detected on one correlation dip, and it is possible that the other measurements contain blended observations."191 1:130DB LD 92855., = HD 92855.192 Revision of the orbit. of TFokovinin (1994).. which= was based on 17 recent measurements. but also on 9 measurements performed between 1916 ancl 1932.," Revision of the orbit of \citet{Toko94}, which was based on 17 recent measurements, but also on 9 measurements performed between 1916 and 1932."193 We applied a correction of |0.381 km + for the former. and |145 km + for the latter.," We applied a correction of +0.381 km $^{-1}$ for the former, and +1.45 km $^{-1}$ for the latter."194 — DD |12 2343, = BD +12 2343.195 AX first orbit was. published. by Jelfries.Bertram&Spurgeon(1995), A first orbit was published by \citet{jeffries95}.196 = HD 102509., = HD 102509.197 Orbit of Crillin&(2004).. partly based. on our. RV measurements.," Orbit of \citet{Griffin04}, partly based on our RV measurements."198 The. periastron epoch was converted in JD. and the systemic velocity was translated in the svsten.," The periastron epoch was converted in JD, and the systemic velocity was translated in the system."199 = BD |21 2357., = BD +21 2357.200 Drift: the RY was decreasing over 10 vears., Drift; the RV was decreasing over 10 years.201 = BD |28 2103., = BD +28 2103.202 Revision of the orbit of(2002): we found a correction of -0.120 kms + to apply to their measurements., Revision of the orbit of; we found a correction of -0.120 km $^{-1}$ to apply to their measurements.203 = LUD 110025., = HD 110025.204 A secondary dip was observed. by Lalbwaehsetal.(2011).. leading to the mass ratio qzz0.64.," A secondary dip was observed by \citet{Halbwachs11}, leading to the mass ratio $q \approx 0.64$."205 = BD |17 2512., = BD +17 2512.206 X secondary dip was observed. by Lalbwachsetal.(2011)..leading to the mass ratio q70.66.," A secondary dip was observed by \citet{Halbwachs11}, leading to the mass ratio $q \approx 0.66$."207 = LD 1yi)6., = HD 110106.208B secondary dip was observed by Lalbwachsetal.120(2011)..m leading to the mass ratio qzzbina0.7!," A secondary dip was observed by \citet{Halbwachs11}, leading to the mass ratio $q \approx 0.75$."209 — ΙΙ) 33., = HD 112033.210 star is ADS 8695. a visual with P=359 2 a=LAS arcsec and Am=2.2 mag (lleintz1991): je secondary component is not. visible on our observations. and the SBI orbit refers to the brightes component of the visual binary.," The star is ADS 8695, a visual binary with $P=359$ yr, $a=1.18$ arcsec and $\Delta m= 2.2$ mag \citep{Heintz97}; the secondary component is not visible on our observations, and the SB1 orbit refers to the brightest component of the visual binary."211 X correction. of -0.263 kms + was applied to the 7 original measurements., A correction of -0.263 km s $^{-1}$ was applied to the 7 original measurements.212 = IUD 234054., = HD 234054.213 A SBI observed over 11 vears. bu with a period still longer.," A SB1 observed over 11 years, but with a period still longer."214 = LD 126661., = HD 126661.215 The RV is slightly decreasing during 6600 davs. until the two dips are separated in our las observations.," The RV is slightly decreasing during 6600 days, until the two dips are separated in our last observations."216 = IID 135117., = HD 135117.217 A SBI observed. over 20 vears. bu with a period still longer.," A SB1 observed over 20 years, but with a period still longer."218 = LD 150631., = HD 150631.219 The variability status is questionable. since £?(v)—1.5'X when the measurement of JL 244993 is discarded.," The variability status is questionable, since $P(\chi^2)=1.5$ when the measurement of JD 2449931 is discarded."220 The BV is then 12.041.2) kms, The RV is then $-12.0 \pm 1.2$ ) km$^{-1}$.221 = LD 153252., = HD 153252.222 X C5-type star without Iuminosity class., A G5-type star without luminosity class.223 Due to the short period. it cannot be a giant: assuming the primary component is a cdwarl. the secondary component has à minimum mass around 50 Jupiter masses. and it is a brown cdwarl canclicla," Due to the short period, it cannot be a giant; assuming the primary component is a dwarf, the secondary component has a minimum mass around 50 Jupiter masses, and it is a brown dwarf candidate."224 = LUD 164025., = HD 164025.225 Orbit of Grillin(2003).. partly based on our RY measurements.," Orbit of \citet{Griffin03}, partly based on our RV measurements."226 The periastron epoch was converted in JD. and the systemic velocity was translated in the system.," The periastron epoch was converted in JD, and the systemic velocity was translated in the system."227 = LD 235865., = HD 238865.228 The star is a triple svstenm. consisting in a long period SBI with an additional short period orbit.," The star is a triple system, consisting in a long period SB1 with an additional short period orbit."229 Preliminary elements of the short. period orbit. were published. by Tokoyinin&Smekhov. (1995): in order to avoid the drift due to the long period. we rejected 21 of our measurements mace before JL 2449000. but we took into account 16 measurements performed with Russian telescopes: the correction to ade to the latter is. |1.15 kim s 2:99," Preliminary elements of the short period orbit were published by \citet{Toko95}; ; in order to avoid the drift due to the long period, we rejected 21 of our measurements made before JD 2449000, but we took into account 16 measurements performed with Russian telescopes; the correction to add to the latter is +1.15 km $^{-1}$."230A= LID 1698522., = HD 169822.231 Revision of the orbit of Lathamct (2002). with a correction of -0.328 km to their measurements.," Revision of the orbit of \citet{Latham02}, with a correction of -0.328 km $^{-1}$ to their measurements."232 “Lhe spectral type of the star is G7 V. leading to à minimum mass around 30 Jupiter masses for the secondary component.," The spectral type of the star is G7 V, leading to a minimum mass around 30 Jupiter masses for the secondary component."233 2:99B= ID 160889., = HD 169889.234 This star was observed as GI41-9 by Lathametal. (2002).. who concluded it had a constant BV.," This star was observed as G141-9 by \citet{Latham02}, , who concluded it had a constant RV."235 A null value of οίv) was obtained from all the LO measurements. and also from the 11 measurements.," A null value of $P(\chi^2)$ was obtained from all the 10 measurements, and also from the 11 measurements."236 Nevertheless. when one outlving measurement is discarded in both sets. £7(v) becomes 0.47 and 0.81. respectively.," Nevertheless, when one outlying measurement is discarded in both sets, $P(\chi^2)$ becomes 0.47 and 0.81, respectively."237 We conclude then that RV is probably constant., We conclude then that the RV is probably constant.238 The RY of the star is then =ihe15.12000wn) d , The RV of the star is then $\bar{V}=(-18.139 \pm 0.114)$ km $^{-1}$ .239= LID 194765., = HD 194765.240 X agonalbright SB2 O nkmmae)with a semi- axis ni lY mas. w should be easily separated.," A bright SB2 (6.70 mag)with a semi-major axis expected around 17 mas, which should be easily separated."241 = ID 214511] =ADS 1611LAD., = HD 214511 =ADS 16111AB.242 Triple system already studied: by Fokovinin.(1998)..., Triple system already studied by \citet{Toko98}. .243 X triple svstem solution was computed. combining a long period SB2 with a 1 as primary component.," A triple system solution was computed, combining a long period SB2 with a SB1 as primary component."244 The period of the SB2was, The period of the SB2was245"Kroupa 2008; Kruijssen Portegies Zwart 2009; Vesperini, McMillan Zwart 2009; 2009), but all point to a shorter Portegiestime to dissolution for Kruijssenloose GCs.","Kroupa 2008; Kruijssen Portegies Zwart 2009; Vesperini, McMillan Portegies Zwart 2009; Kruijssen 2009), but they all point to a shorter time to dissolution for loose GCs."246" For theythis reason, it is important to cover both dense (c> 1.5) and loose (c« 1.5) clusters when studying the temporal evolution of the stellar MF."," For this reason, it is important to cover both dense $c > 1.5$ ) and loose $c< 1.5$ ) clusters when studying the temporal evolution of the stellar MF."247" As regards YCs, the data cover high quality observations of fourteen star-forming regions and associations of various ages, as shown in reftabl.."," As regards YCs, the data cover high quality observations of fourteen star-forming regions and associations of various ages, as shown in \\ref{tab1}."248" For each YC, we searched for the most accurate determination of the MF in the recent literature (see references in reftab1)), specifically for information on the number of stars measured lookingin each mass bin."," For each YC, we searched for the most accurate determination of the MF in the recent literature (see references in \\ref{tab1}) ), looking specifically for information on the number of stars measured in each mass bin."249" In those cases in which the MFs were not directly available in tabular form, we extracted the data points (i.e. number of stars in each mass bin) from the published graphs, but we specifically ignored any fits (with power-law or log-normal functions) or interpolations of these data that the various authors may have carried out."," In those cases in which the MFs were not directly available in tabular form, we extracted the data points (i.e. number of stars in each mass bin) from the published graphs, but we specifically ignored any fits (with power-law or log-normal functions) or interpolations of these data that the various authors may have carried out."250 The MFs in our sample were determined either by converting an observed luminosity function via a theoretical mass—luminosity relationship or by counting the number of stars falling between evolutionary tracks in a color-magnitude or Hertzsprung-Russell diagram., The MFs in our sample were determined either by converting an observed luminosity function via a theoretical mass--luminosity relationship or by counting the number of stars falling between evolutionary tracks in a color–magnitude or Hertzsprung–Russell diagram.251" We only restricted our search to works in which the MF is based on known cluster members and a correction for photometric has been applied to the luminosity function, ifincompleteness needed."," We only restricted our search to works in which the MF is based on known cluster members and a correction for photometric incompleteness has been applied to the luminosity function, if needed."252" Furthermore, to guarantee that the sample is as homogeneous as possible, we have ignored any corrections to the MF meant to account for the presence of binaries, since these are rather uncertain and only available in a limited number of cases."," Furthermore, to guarantee that the sample is as homogeneous as possible, we have ignored any corrections to the MF meant to account for the presence of binaries, since these are rather uncertain and only available in a limited number of cases."253" Therefore, the MFs that we consider in this study, for both YCs and GCs, are the MFs of stellar systems, i.e. we do not distinguish between single and multiple stars."," Therefore, the MFs that we consider in this study, for both YCs and GCs, are the MFs of stellar systems, i.e. we do not distinguish between single and multiple stars."254" In order to prevent biases in the results, the data should refer to the global MF (GMP), i.e. to the MF of the cluster as a whole."," In order to prevent biases in the results, the data should refer to the global MF (GMF), i.e. to the MF of the cluster as a whole."255" For GCs, where complete cluster coverage is not always possible, we used information on mass stratification and mass segregation to derive the GMF from the local MF (see e.g. De Marchi et al."," For GCs, where complete cluster coverage is not always possible, we used information on mass stratification and mass segregation to derive the GMF from the local MF (see e.g. De Marchi et al."256 2006)., 2006).257" Alternatively, we used the MF measured near the half-light radius (i.e. the effective radius containing half of the cluster's luminosity, see e.g. Portegies Zwart, McMillan Gieles 2010), since it has been shown to reflect quite reliably the properties of the GMF (De Marchi et al."," Alternatively, we used the MF measured near the half-light radius (i.e. the effective radius containing half of the cluster's luminosity, see e.g. Portegies Zwart, McMillan Gieles 2010), since it has been shown to reflect quite reliably the properties of the GMF (De Marchi et al."258 2000)., 2000).259" For YCs, we specifically selected those objects and studies for which the coverage is as complete as possible."," For YCs, we specifically selected those objects and studies for which the coverage is as complete as possible."260" We then performed a multivariate fit to each MF with a TPL distribution and derived the values of the parameters o, B and m, that simultaneously provide the smallest residuals."," We then performed a multivariate fit to each MF with a TPL distribution and derived the values of the parameters $\alpha$, $\beta$ and $m_c$ that simultaneously provide the smallest residuals."261 The values of the best fitting parameters and their associated 1c uncertainties are given in reftabl for all the clusters in our sample.," The values of the best fitting parameters and their associated $1\,\sigma$ uncertainties are given in \\ref{tab1} for all the clusters in our sample."262" The MFs of a few selected clusters are also shown graphically in reffig1, together with the best TPL fits as per the parameters of reftabl.."," The MFs of a few selected clusters are also shown graphically in \\ref{fig1}, together with the best TPL fits as per the parameters of \\ref{tab1}."263" We compare and discuss the results of our best fits in 55, but we precede that with a short discussion on the value of a."," We compare and discuss the results of our best fits in 5, but we precede that with a short discussion on the value of $\alpha$."264" In the TPL distribution described 11, the value of α determines the shape of the MF byabove Equationthe characteristic mass m,."," In the TPL distribution described by 1, the value of $\alpha$ determines the shape of the MF above the characteristic mass $m_c$."265" Since the data for GCs do not constrain the mass range above the main-sequence turn-off (> 0.8MMo), we had to assume a value of a for those objects."," Since the data for GCs do not constrain the mass range above the main-sequence turn-off $> 0.8$ $_\odot$ ), we had to assume a value of $\alpha$ for those objects."266" An obvious first choice would be the slope of the Salpeter (1955) IMF, namely à;=—2.35."," An obvious first choice would be the slope of the Salpeter (1955) IMF, namely $\alpha=-2.35$."267" According to Kroupa (2002), that is the average value of the slope of the stellar MF in clusters and associations above 0.5Msolar."," According to Kroupa (2002), that is the average value of the slope of the stellar MF in clusters and associations above $0.5$."268". However, this is not a meaningful average because of the very large scatter at m>1 displayed by the MF slopes included in Kroupa’s (2002) compilation, which in turn are taken from Scalo (1998)."," However, this is not a meaningful average because of the very large scatter at $ m>2691$ displayed by the MF slopes included in Kroupa's (2002) compilation, which in turn are taken from Scalo (1998)."270" In fact, after discussing with this latter author (J. Scalo, priv."," In fact, after discussing with this latter author (J. Scalo, priv."271" comm.),"," comm.),"272 we realised that the o=—2.35 value obtained in that way is not a average., we realised that the $\alpha=-2.35$ value obtained in that way is not a meaningful average.273" It is also not consistent with any of the individual o meaningfulvalues that we measure for the YCs in reftabl,, except for Taurus and possibly Pleiades."," It is also not consistent with any of the individual $\alpha$ values that we measure for the YCs in \\ref{tab1}, , except for Taurus and possibly Pleiades."274substantially..,.275citepjorgenssen2004a.. while the eritical density of iis ~107 citepho1977..,", while the critical density of is $\sim10^4$ \\citep{ho1977}."276 From the dust continuum emission we estimated a density of ~10° im the central core., From the dust continuum emission we estimated a density of $\sim10^6$ in the central core.277" In addition. the density obtained from single-dish observations at 850 pm (Jennessetal. 1995).. adopting a core size of ~26"", is ~4»10°em™.."," In addition, the density obtained from single-dish observations at 850 $\mu$ m \citep{Jenness1995}, , adopting a core size of $\sim26''$, is $\sim4\times10^5$."278 For these densities the transitions of both aand aare thermalized. and thus excitation effects do not seem appropriate to explain the eemission drop in the central core.," For these densities the transitions of both and are thermalized, and thus excitation effects do not seem appropriate to explain the emission drop in the central core."279 On the other hand. the hyperfine structure method in CLASS assumes equal excitation temperature for all the hyperfine components.," On the other hand, the hyperfine structure method in CLASS assumes equal excitation temperature for all the hyperfine components."280 Danieletal.(2006) show that this assumption does not hold for the case of high opacities due to radiative processes., \citet{daniel2006} show that this assumption does not hold for the case of high opacities due to radiative processes.281" Since in 55142 we derived an opacity of Tyo, =0.3-0.6 for the mmolecule. the problem reported by Danieletal.(2006) is not likely affecting our result."," Since in 5142 we derived an opacity of $\tau_{\mathrm{TOT}}\simeq$ 0.3–0.6 for the molecule, the problem reported by \citet{daniel2006} is not likely affecting our result."282 The high values. up to 1000. of the aabundance ratio found in the central core of 55142 can be reproduced by our chemical model for high densities (1x10° em4) and high temperaures (T.=~70 ΚΙ.," The high values, up to 1000, of the abundance ratio found in the central core of 5142 can be reproduced by our chemical model for high densities $n\simeq10^6$ ) and high temperaures $T\simeq70$ K)."283 The chemical modeling performed in Sect., The chemical modeling performed in Sect.284 5 indicates that both density and temperature play an important role in determining the molecular abundance of aand and hence their ratio.," 5 indicates that both density and temperature play an important role in determining the molecular abundance of and, and hence their ratio."285 The central core has a higher temperature and density than the western and eastern cores., The central core has a higher temperature and density than the western and eastern cores.286 Thus. a different chemistry can develop due to CO evaporation from the grain mantles.," Thus, a different chemistry can develop due to CO evaporation from the grain mantles."287 The CO desorption in the central core of 55142 leads to the destruction ofN->H7.. and hence the aabundance ratio increases considerably relative to the value found in the westert and eastern cores.," The CO desorption in the central core of 5142 leads to the destruction of, and hence the abundance ratio increases considerably relative to the value found in the western and eastern cores."288 This is supported by the fact that iis not frozen out in the central core. whereas it is faintly detected in the western and eastern cores (Zhangetal.2007).," This is supported by the fact that is not frozen out in the central core, whereas it is faintly detected in the western and eastern cores \citep{zhang2007}."289. As pointed out in. Sect 5.2.2. Collingsetal.(2004). find experimental evidences that the desorption of ffrom grain mantles takes place at a temperature of ~120 K. Therefore. since m our model we assumed a maximum temperature of 70 K. the high aabundance ratio in the central core is mainly a consequence of the destruction of bby CO rather than an enhancement of iin the gas phase.," As pointed out in Sect 5.2.2, \citet{collings2004} find experimental evidences that the desorption of from grain mantles takes place at a temperature of $\sim120$ K. Therefore, since in our model we assumed a maximum temperature of 70 K, the high abundance ratio in the central core is mainly a consequence of the destruction of by CO rather than an enhancement of in the gas phase."290 The disappearance of ffrom the gas phase has been reported by several authors. both in low- and high-mass star-forming regions.," The disappearance of from the gas phase has been reported by several authors, both in low- and high-mass star-forming regions."291 In the low-mass regime. there are several reports of central ddepletion in Class 0/Class I protostars 00419141522: Belloche&André2004:: 11623: DiFrancescoetal.2004: L483 mm: Jorgensen2004: Barnard Ie: Matthewsetal.2006 BHR HIIRSI: al.2008)).," In the low-mass regime, there are several reports of central depletion in Class 0/Class I protostars 04191+1522: \citealt{belloche2004}; 1623: \citealt{diFrancesco2004}; L483 mm: \citealt{jorgenssen2004a}; Barnard 1c: \citealt{Matthews2006}; BHR IRS1: \citealt{chen2008}) )."292" Regarding the high-mass regime. Pirogovetal.(2003.2007) observed with a single-dish telescope à sample of dense cores associated with massive stars and star clusters containing IRAS point sources. and find that. for most of the sources, there is a decrease of the aabundance toward the dust column density peak."," Regarding the high-mass regime, \citet{Pirogov2003,Pirogov2007} observed with a single-dish telescope a sample of dense cores associated with massive stars and star clusters containing IRAS point sources, and find that, for most of the sources, there is a decrease of the abundance toward the dust column density peak."293 Furthermore. interferometric observations of ttoward the high-mass star-forming regions 2230334595] (Reid&Matthews2008) also reveal a significant destruction of.," Furthermore, interferometric observations of toward the high-mass star-forming regions 23033+5951 \citep{Reid2008} also reveal a significant destruction of."294. Therefore. the high ratio measured toward the central core of 55142 seems to be due to the rapid destruction of wwhen CO ts released from grain mantles as seems to be the case for the other regions cited above.," Therefore, the high ratio measured toward the central core of 5142 seems to be due to the rapid destruction of when CO is released from grain mantles as seems to be the case for the other regions cited above."295 Although in our chemical model we did not include molecular outflows. they may strongly affect the surrounding dense gas. and hence the aabundance ratio.," Although in our chemical model we did not include molecular outflows, they may strongly affect the surrounding dense gas, and hence the abundance ratio."296 As already discussed by Chenetal.(2008).. who propose three stages of the interaction between the eemission and molecular outflows. based on the morphology of the ddense gas and the jet/outflow emission.," As already discussed by \citet{chen2008}, who propose three stages of the interaction between the emission and molecular outflows, based on the morphology of the dense gas and the jet/outflow emission."297 Here we suggests a qualitative picture of the evolution of the aabundance ratio., Here we suggests a qualitative picture of the evolution of the abundance ratio.298 First. in starless cores with high densities most molecules are highly depleted. favoring the formation of aandN>H™.," First, in starless cores with high densities most molecules are highly depleted, favoring the formation of and."299. It seems that the CO depletionfavors the formation of against in core centers at high densities., It seems that the CO depletionfavors the formation of against in core centers at high densities.300 Thus. the rratio reaches high values. consistent with the results found in low-mass starless cores (Casellietal.2002a:Hotzelal.2004;Friesenet2010) as well as for starless cores nearby massive stars. like in the high-mass star-forming regions 22029343952 (Palauetal.2007).," Thus, the ratio reaches high values, consistent with the results found in low-mass starless cores \citep{caselli2002a,hotzel2004,friesen2010} as well as for starless cores nearby massive stars, like in the high-mass star-forming regions 20293+3952 \citep{palau2007}."301. The high rratio 1n. starless cores has been reproduced by our chemical model shown in Fig.9 (left panels)., The high ratio in starless cores has been reproduced by our chemical model shown in \ref{modelfigs}~ (left panels).302 After this phase. once the star has formed. the rratio decreases. as found in the low-mass regime (Casellietal.2002a:Hotzeletal.2004;Friesen 2010).. and in YSOs close to the massive star 22029343952: Palauetal.2007:: 000117+612: Busquet iin prep.).," After this phase, once the star has formed, the ratio decreases, as found in the low-mass regime \citep{caselli2002a,hotzel2004,friesen2010}, and in YSOs close to the massive star 20293+3952: \citealt{palau2007}; 00117+612: Busquet in prep.)."303 In this situation. CO is less depleted. and mmolecules inthe envelope are entrained by the molecular outflow. being the eemission elongated in the direction of the molecular outflow. as found in 2202933952. (Palauetal.2007) and 00011746412 (Palauetal. 2010)..," In this situation, CO is less depleted, and molecules inthe envelope are entrained by the molecular outflow, being the emission elongated in the direction of the molecular outflow, as found in 20293+3952 \citep{palau2007} and 00117+6412 \citep{palau2010}. ."304 Finally. the next," Finally, the next"305limit.,limit.306 So. we can assume that both effects cancel out and 2.2 mG is a [air estimation for the POS field strength.," So, we can assume that both effects cancel out and 2.2 mG is a fair estimation for the POS field strength."307 The mass-to-flux ratio gives information on whether the magnetic field can support the cloud against (he eravilational collapse and. therefore. it provides clues about. the evolutionary state of the source.," The mass-to-flux ratio gives information on whether the magnetic field can support the cloud against the gravitational collapse and, therefore, it provides clues about the evolutionary state of the source."308" Specifically, (his quantity compares (he pressure produced bv an amount of mass M in a magnelic tube of [lux e."," Specifically, this quantity compares the pressure produced by an amount of mass M in a magnetic tube of flux $\Phi$."309 A critical value. reachecl when {he magnetic pressure is no longer able to support the gravitational pulling. is given by (2*328VG)vC)+ (Nakano(Nak:&NakamuraNaki‘a19178).," A critical value, reached when the magnetic pressure is no longer able to support the gravitational pulling, is given by $(2\pi\sqrt{G})^{-1}$ \citep{Nakano78}."310.1978).. Observationallv.Observationally. this parameterpari|ds is defined.defined bsby (Cruteherοἱal.1999): where (1Φα 1s (e mass-to-fhix ratio of an uniform disk where eravily is balanced bv magnetic pressure. =2.8:»j allowing for He. A is the cloud area covered by observations. N(Ils) is the column density and 2 is the magnetic field strength.," Observationally, this parameter is defined by \citep{Crutcher99}: where $(M/\Phi)_{critical}$ is the mass-to-flux ratio of an uniform disk where gravity is balanced by magnetic pressure, $m = 2.8m_{H}$ allowing for He, $A$ is the cloud area covered by observations, $_{2}$ ) is the column density and $B$ is the magnetic field strength."311 Applying ihe POS magnetic field strength obtained in the previous paragraph. 5=2.2 mG. and the column densities derived in section 3.1.. we estimate a mass-to-[Iux ratio for FIR. 5À of 1.6 (for --60 IX). which corresponds to a core in a supercritical stage.," Applying the POS magnetic field strength obtained in the previous paragraph, $B = 2.2$ mG, and the column densities derived in section \ref{prop_dust}, we estimate a mass-to-flux ratio for FIR 5A of 1.6 (for $T = 60$ K), which corresponds to a core in a supercritical stage."312" Ir anv case. (hose ealeulations are restricted to the dust envelope. without taking into account the mass contribution of the embedded. protostar,"," In any case, those calculations are restricted to the dust envelope, without taking into account the mass contribution of the embedded protostar."313 We consider that the derived mass-to-Ilux. values are only a lower limit for (his quantity and therefore it is in agreement. wilh the observed star-Dorming signatures., We consider that the derived mass-to-flux values are only a lower limit for this quantity and therefore it is in agreement with the observed star-forming signatures.314 We can also derive the ratio between (turbulent. ancl magnetic energies., We can also derive the ratio between turbulent and magnetic energies.315 From the autocorrelation Iunction of the polarization position angles. il is possible to measure how the dispersion of PAs varies with respect to the distinct length scales within the cloud.," From the autocorrelation function of the polarization position angles, it is possible to measure how the dispersion of PA's varies with respect to the distinct length scales within the cloud."316 This [function provides an indirect calculation of the turbulent to magnetic energy as, This function provides an indirect calculation of the turbulent to magnetic energy as317up to 10 arcmin.,up to 10 arcmin.318" Comparing these results with the size of the error bars shown in the upper panel, we can conclude that with relatively small maps, like those analyzed in this paper, we can use the skewness statistics to detect only strong primordial non-Gaussianities (funn=+500, "," Comparing these results with the size of the error bars shown in the upper panel, we can conclude that with relatively small maps, like those analyzed in this paper, we can use the skewness statistics to detect only strong primordial non-Gaussianities $f_{\rm NL}=\pm 500, \pm 1000$ )."319"As shown by the results of the analysis made by |Fedeliwouldwe&Moscardini| only with data covering very large area it be possible to (2010),,disentangle the effects of non-Gaussian models with fwr,=+100."," As shown by the results of the analysis made by \cite{Fedeli2010}, only with data covering very large area it would be possible to disentangle the effects of non-Gaussian models with $f_{\rm NL}=\pm 100$."320" Finally we notice that a similar analysis performed on the shear field, which is directly related to the convergence one, would provide comparable results."," Finally we notice that a similar analysis performed on the shear field, which is directly related to the convergence one, would provide comparable results."321" An important theoretical quantity, directly related to the matter power spectrum, is the effective convergence power spectrum P,.(£), which is defined as the squared modulus of the Fourier transform of «, averaged on the modes having a given multipole 6."," An important theoretical quantity, directly related to the matter power spectrum, is the effective convergence power spectrum $P_{\kappa}(\ell)$, which is defined as the squared modulus of the Fourier transform of $\kappa$, averaged on the modes having a given multipole $\ell$."322" Starting from P.(6). it is possible to derive analytic expressions for the power spectra of the other lensing quantities, such as those for the shear (P,) and the two flexions (Pr and Pc), namely: In Fig."," Starting from $P_{\kappa}(\ell)$, it is possible to derive analytic expressions for the power spectra of the other lensing quantities, such as those for the shear $P_{\gamma}$ ) and the two flexions $P_{F}$ and $P_{G}$ ), namely: In Fig."323" we show, for the Gaussian model only, the comparison betweenB] the power spectra extracted from our simulations and the corresponding theoretical predictions including the non-linear effect."," \ref{fig:PS} we show, for the Gaussian model only, the comparison between the power spectra extracted from our simulations and the corresponding theoretical predictions including the non-linear effect."324" We consider different lensing quantities: effective convergence, shear and reduced shear g=y/(1—κ) (upper panel) and two components of the flexion (bottom panel)."," We consider different lensing quantities: effective convergence, shear and reduced shear $g\equiv\gamma/(1-\kappa)$ (upper panel) and two components of the flexion (bottom panel)."325" We notice that the agreement between the numerical results and the theoretical expectations holds, as expected, only in a given range of wavenumbers."," We notice that the agreement between the numerical results and the theoretical expectations holds, as expected, only in a given range of wavenumbers."326" On large scales (small @), the departure from theory is due to the low number of modes we can use to average the numerical power spectra."," On large scales (small $\ell$ ), the departure from theory is due to the low number of modes we can use to average the numerical power spectra."327" On small scales (large £), the disagreement is produced by the mass resolution (particularly relevant for the flexion power spectrum) and by numerical artifacts produced by the weighting function appearing in the definition of a given power spectrum."," On small scales (large $\ell$ ), the disagreement is produced by the mass resolution (particularly relevant for the flexion power spectrum) and by numerical artifacts produced by the weighting function appearing in the definition of a given power spectrum."328 We can use Fig., We can use Fig.329 as a reference to define the range of £ on which the numerical results can be trusted: typically 10?€£10%., \ref{fig:PS} as a reference to define the range of $\ell$ on which the numerical results can be trusted: typically $ 10^2 \la \ell\la 10^4$.330 In Fig., In Fig.331 B] we show the ratio between the power spectra derived from the non-Gaussian and the Gaussian simulations., \ref{fig:PSratio} we show the ratio between the power spectra derived from the non-Gaussian and the Gaussian simulations.332" We present here only the results for the effective convergence since the other power spectra can be obtained by introducing the suitable dependence on £, and then the corresponding ratios are identical."," We present here only the results for the effective convergence since the other power spectra can be obtained by introducing the suitable dependence on $\ell$, and then the corresponding ratios are identical."333" As expected, the effects of the primordial non-Gaussianity become more evident when the absolute value of the parameter fwr increases."," As expected, the effects of the primordial non-Gaussianity become more evident when the absolute value of the parameter $f_{\rm NL}$ increases."334" At very large scales, where the evolution is still in the linear regime, the Gaussian model and all the non-Gaussian ones here considered display the same power spectrum, and the ratio is in practice unity."," At very large scales, where the evolution is still in the linear regime, the Gaussian model and all the non-Gaussian ones here considered display the same power spectrum, and the ratio is in practice unity."335 At smaller scales ({>100— 200) the behaviors, At smaller scales $\ell>100-200$ ) the behaviors336"microgravity experiments, this mechanism is connected to a second stepper motor, which turns the target by 180° shortly before the drop capsule is decelerated.","microgravity experiments, this mechanism is connected to a second stepper motor, which turns the target by $180^\circ$ shortly before the drop capsule is decelerated."337" Once turned, a second lifting solenoid lowers a cap over the target body to prevent mass loss."," Once turned, a second lifting solenoid lowers a cap over the target body to prevent mass loss."338" The impacts are recorded with a high-speed camera with a frame rate of 2,000 Hz and a field of view of 2x1cm?."," The impacts are recorded with a high-speed camera with a frame rate of 2,000 Hz and a field of view of $2 \times 1\ \mathrm{cm}^2$."339" For the back- illumination of the field of view, we used a 150 W halogen lamp and a diffusor."," For the back-light illumination of the field of view, we used a 150 W halogen lamp and a diffusor."340" As analog material for the protoplanetary dust, we used for both, the target and projectile aggregates, monodisperse, spherical SiOs particles with a diameter of 1.5 ym. The properties of these grains are well known from measurements of ?,, aggregate properties have been measured by ? and in Paper I, and it was also used in many of the experiments reviewed by ?.."," As analog material for the protoplanetary dust, we used for both, the target and projectile aggregates, monodisperse, spherical $_2$ particles with a diameter of 1.5 $\mu$ m. The properties of these grains are well known from measurements of \citet{HeimEtal:1999}, , aggregate properties have been measured by \citet{BlumSchraepler:2004} and in Paper I, and it was also used in many of the experiments reviewed by \citet{BlumWurm:2008}."341" The spherical shape of the dust monomers allows a direct comparison to numerical modeling of dust aggregates (e.g.12) and the same material was also used for the calibration of an SPH code (PaperIV,"," The spherical shape of the dust monomers allows a direct comparison to numerical modeling of dust aggregates \citep[e.g.][]{DominikTielens:1997, WadaEtal:2009} and the same material was also used for the calibration of an SPH code \citep[Paper IV,][]{GeretshauserEtal:2010}."342" We used 1.5 mm sized ?)..fragments of highly porous dust agglomerates with a volume filling factor of $9=0.15+0.01 formed by the method described in Paper I. The fragments were cut out of the larger agglomerate with a razor blade, which causes only a minor compaction at the edge of the agglomerates."," We used 1.5 mm sized fragments of highly porous dust agglomerates with a volume filling factor of $\phi_0 = 0.15 \pm 0.01$ formed by the method described in Paper I. The fragments were cut out of the larger agglomerate with a razor blade, which causes only a minor compaction at the edge of the agglomerates."343" Before each experiment, 25 of these fragments were placed in the magazine."," Before each experiment, 25 of these fragments were placed in the magazine."344 'The dust-aggregate target is cylindrical in shape with 1 cm diameter and thickness of a few mm and is made of compacted and sintereda dust., The dust-aggregate target is cylindrical in shape with 1 cm diameter and a thickness of a few mm and is made of compacted and sintered dust.345" The SiOs dust particles are compacted with forces from 0.6 to 3 kN, which yield a volume filling factor of about 0.4 to 0.5 (Paper IV)."," The $_2$ dust particles are compacted with forces from 0.6 to 3 kN, which yield a volume filling factor of about 0.4 to 0.5 (Paper IV)."346 These pellets were then sintered for one hour at 1100° C. The sintering prevents the target from damage by the impacts and thus provides an indestructible target from the same material., These pellets were then sintered for one hour at $1100^\circ$ C. The sintering prevents the target from damage by the impacts and thus provides an indestructible target from the same material.347 The surface structure of the target is not changed by the sintering process and is corresponding to a plain layer of dust and allows a self-consistent transition from the indestructible target to the self consistently grown dust layer., The surface structure of the target is not changed by the sintering process and is corresponding to a plain layer of dust and allows a self-consistent transition from the indestructible target to the self consistently grown dust layer.348" Before each experimental sequence, the target was weighted (see below) and placed in the target holding mechanism."," Before each experimental sequence, the target was weighted (see below) and placed in the target holding mechanism."349 The mass of the target was determined after each sequence of 25 shots in the laboratory and after each series typically a little less than 25 in the microgravity(of experiment., The mass of the target was determined after each sequence of 25 shots in the laboratory and after each series (of typically a little less than 25 shots) in the microgravity experiment.350" Since the dust shots)adsorbs air humidity, the determination of the mass and, thus, the calculation of the volume filling factor and the accretion efficiency must be conducted with great care."," Since the dust adsorbs air humidity, the determination of the mass and, thus, the calculation of the volume filling factor and the accretion efficiency must be conducted with great care."351 First experiments showed that the measured mass of larger amounts of dust fluctuates between different weighings and moreover grows on average with the time after the aeration., First experiments showed that the measured mass of larger amounts of dust fluctuates between different weighings and moreover grows on average with the time after the aeration.352" To estimate the error caused by this effect, we used the standard deviation of five different weighings within a short time."," To estimate the error caused by this effect, we used the standard deviation of five different weighings within a short time."353 This error is typically of the order 0.2 mg and was found to be independent of the dust mass in the relevant range., This error is typically of the order 0.2 mg and was found to be independent of the dust mass in the relevant range.354" Thus, we regard this error as a limitation to the accuracy of the balance with respect to its place of installation."," Thus, we regard this error as a limitation to the accuracy of the balance with respect to its place of installation."355" For the drop tower experiments and early laboratory series (series 1-g, 3a-lab, and 3b-lab, cf."," For the drop tower experiments and early laboratory series (series $\mu$ g, 3a-lab, and 3b-lab, cf."356" Table Π)), we used this mass determination with a time delay of less than 5 minutes between aeration and weighing."," Table \ref{tab.exp_overview}) ), we used this mass determination with a time delay of less than 5 minutes between aeration and weighing."357" For the later laboratory experiments (series 2-lab and 4-lab), we used an air-tight weighing container to fix the content of water."," For the later laboratory experiments (series 2-lab and 4-lab), we used an air-tight weighing container to fix the content of water."358" Continuous measurements over half an hour showed that the mass also grows during the monitored weighing process, which we attribute to (slowly diminishing) surface charges on the glass container."," Continuous measurements over half an hour showed that the mass also grows during the monitored weighing process, which we attribute to (slowly diminishing) surface charges on the glass container."359" The temporal growth in mass can be approximated by an exponential decay function where is the measured target mass at time { after the aeration M(t)and Am and 7 are fitting constants of the order of 0.5 to 1.5 mg and 4 to 12 min, respectively."," The temporal growth in mass can be approximated by an exponential decay function where $M(t)$ is the measured target mass at time $t$ after the aeration and $\Delta m$ and $\tau$ are fitting constants of the order of 0.5 to 1.5 mg and 4 to 12 min, respectively."360" The value Myeal, to which Equation (1) converges, is our mass determination for which we still use an error of 0.2 mg mentioned above to consider the accuracy of the balance."," The value $M_\mathrm{real}$, to which Equation \ref{Eq._Mass_det}) ) converges, is our mass determination for which we still use an error of 0.2 mg mentioned above to consider the accuracy of the balance."361" To reduce the influence of variations in the air humidity (ie. water content in the container), we furthermore used the differences between two weighings before and after each experimental sequence for further evaluations."," To reduce the influence of variations in the air humidity (i.e. water content in the container), we furthermore used the differences between two weighings before and after each experimental sequence for further evaluations."362" With respect to the mechanical sensitivity of the porous projectiles, it is not feasible to determine the mass of every single projectile."," With respect to the mechanical sensitivity of the porous projectiles, it is not feasible to determine the mass of every single projectile."363" Likewise, it is not possible to weigh the whole magazine due to its largesize and mass."," Likewise, it is not possible to weigh the whole magazine due to its largesize and mass."364 We therefore used an average projectile, We therefore used an average projectile365bands.,bands.366" For canonical parameters Vi.)=300 keV. αι=I. and às= 2.3). the ratio vfiaic/if.m at 100 MeV is about 0.01(εερ)>,/400)77, assuming that the synchrotron spectrum continues into the GeV window without a break and IC is not much suppressed by the Klein-Nishina effect."," For canonical parameters $h \nu_{\rm syn} = 300$ keV, $\alpha_1 = 1$ , and $\alpha_2 = 2.3$ ), the ratio $\nu f_{\nu,{\rm IC}} / \nu367f_{\nu,{\rm syn}}$ at 100 MeV is about $0.01 (\epsilon_{e} /368\epsilon_B)^{1/2} (\gamma_m / 400)^{-2}$, assuming that the synchrotron spectrum continues into the GeV window without a break and IC is not much suppressed by the Klein-Nishina effect."369 Therefore. the synchrotron component dominates around the lower-energy limit where most of the photons (although not most of the fluence) are observed.," Therefore, the synchrotron component dominates around the lower-energy limit where most of the photons (although not most of the fluence) are observed."370 In the case of EGRET. since only a handful of photons were detected in all EGRET events. these are expected to be dominated by the synchrotron low-energy (~100 MeV) photons.," In the case of EGRET, since only a handful of photons were detected in all EGRET events, these are expected to be dominated by the synchrotron low-energy $\sim$ 100 MeV) photons."371 This indicates that the quantity we can constrain using the EGRET fluence upper limits is not ic but is=οί100ΜΟΝ)Foy. the ratio of synchrotror fluence around 100 MeV and that in the MeV range.," This indicates that the quantity we can constrain using the EGRET fluence upper limits is not $\eta_{\rm IC}$ but $\eta_{\rm syn} = F_{\rm syn}(100 ~\mathrm{MeV}) / F_{\rm MeV}$, the ratio of synchrotron fluence around 100 MeV and that in the MeV range."372 li addition. this picture is indeed consistent with the fact that the spectral indices of GeV photons for several GRBs measurec with EGRET are à= 2-31994)..," In addition, this picture is indeed consistent with the fact that the spectral indices of GeV photons for several GRBs measured with EGRET are $\alpha = 2$ –3."373 Note however. that the in GLAST-LAT and EGRET bands can be dominatec by a much harder IC component (4z 1-2) that peaks above ] GeV and may carry up to ~10 times more energy thar the one observed at 100 MeV without being detected.," Note however, that the in -LAT and EGRET bands can be dominated by a much harder IC component $\alpha \approx 1$ –2) that peaks above $\sim$ 1 GeV and may carry up to $\sim$ 10 times more energy than the one observed at 100 MeV without being detected."374 This is because even when the ~10 GeV fluence is ten times larger. the small photon number at such high-energies ts still small enough to avoid detection.," This is because even when the $\sim$ 10 GeV fluence is ten times larger, the small photon number at such high-energies is still small enough to avoid detection."375 Thus. EGRET observations. which are consistent with measurement of the synchrotron high energy tail. can only put an upper limit on ijc.," Thus, EGRET observations, which are consistent with measurement of the synchrotron high energy tail, can only put an upper limit on $\eta_{\rm IC}$."376 The afterglow is considered to be a synchrotron emission from electrons accelerated in the external shock. which ts caused by the interaction between the relativistic ejecta and the interstellar medium.," The afterglow is considered to be a synchrotron emission from electrons accelerated in the external shock, which is caused by the interaction between the relativistic ejecta and the interstellar medium."377 In this model. the synchrotron emission dominates the spectrum from radio to X-ray.," In this model, the synchrotron emission dominates the spectrum from radio to X-ray."378 The associated IC. emission is expected to dominate the GeV energy range (Le.. ic> i4). Since the electron Lorentz factor is much larger than the case of prompt emission (see eq. [10]].," The associated IC emission is expected to dominate the GeV energy range (i.e., $\eta_{\rm IC} \gg \eta_{\rm379syn}$ ), since the electron Lorentz factor is much larger than the case of prompt emission (see eq. \ref{eq:gamma_e}] ],"380" where the relative and bulk Lorentz factors are the same. L4= L4). compensating the smaller 14,44 (eq. [3]])."," where the relative and bulk Lorentz factors are the same, $\Gamma_{\rm rel} = \Gamma_{b}$ ), compensating the smaller $\nu_{\rm syn}$ (eq. \ref{eq:frequency}] ])."381" During the first several minutes (observer time). electrons might be cooling fast (n4= 1.5) with fy,z| keV. while 7,7 107-10."," During the first several minutes (observer time), electrons might be cooling fast $\alpha_1 = 1.5$ ) with $h\nu_{\rm syn} \approx 1$ keV, while $\gamma_{m} \approx38210^4$ $^5$."383" This implies that the fraction of the IC energy that falls in GLAST-LAT energy window is close to unity. Le. &,2: 0.2-0.9 from equation (7)) (for EGRET 4,0.08-- 0.5) and «knz 0.7-1 from equations (5))-(6))."," This implies that the fraction of the IC energy that falls in -LAT energy window is close to unity, i.e, $\xi_{w} \approx 0.2$ –0.9 from equation \ref{eq:eta_w 1}) ) (for EGRET $\eta_{w} \approx 0.08$ --0.5) and $\xi_{\rm KN} \approx 0.7$ –1 from equations\ref{eq:eta KN}) \ref{eq:gamma KN}) )."384 Since ης at early time Is close to the upper limit of the energy window the effective photon index of the IC emission within the detector window during this time is z1.5-2., Since $h\nu_{\rm IC}$ at early time is close to the upper limit of the energy window the effective photon index of the IC emission within the detector window during this time is $\approx$ 1.5–2.385" At later times the electrons are at the slow-cooling regime and 744 18 the cooling frequency. while a typical , is the Lorentz factor of electrons that cooled significantly2001)."," At later times the electrons are at the slow-cooling regime and $\nu_{\rm syn}$ is the cooling frequency, while a typical $\gamma_{e}$ is the Lorentz factor of electrons that cooled significantly."386. In this regime the SSC peak is very broad and its location is almost constant with time., In this regime the SSC peak is very broad and its location is almost constant with time.387 For typical parameters. the Klein-Nishina effect do not play a major role while the peak of the SSC emission falls within GLAST-LAT and EGRET windows.," For typical parameters, the Klein-Nishina effect do not play a major role while the peak of the SSC emission falls within -LAT and EGRET windows."388" Therefore. at late time £,~| and the effective photon index within the energy windows of these detectors 1s 222."," Therefore, at late time $\xi_w \approx 1$ and the effective photon index within the energy windows of these detectors is $\approx$ 2."389 One should. however. note that on long time scales the GeV background becomes important. making it hard to detect the GeV afterglow.," One should, however, note that on long time scales the GeV background becomes important, making it hard to detect the GeV afterglow."390" Therefore. the optimal time scale for GeV afterglow search would be ~ 100-10"" s2001).. The afterglow GeV fluence. Foey in equation (1)). is that integrated over a given time scale. while [μον is collected over roughly Του. during which of the MeV photons are counted."," Therefore, the optimal time scale for GeV afterglow search would be $\sim$ $^3$ s. The afterglow GeV fluence, $F_{\rm GeV}$ in equation \ref{eq:EGRET fluence}) ), is that integrated over a given time scale, while $F_{\rm MeV}$ is collected over roughly $T_{90}$, during which of the MeV photons are counted."391 The total energy radiated away by the radio to ray afterglow during every decade of time is roughly 0.01—0.1 of the energy emitted in the prompt phase., The total energy radiated away by the radio to X-ray afterglow during every decade of time is roughly 0.01--0.1 of the energy emitted in the prompt phase.392" Therefore we expect a bright GeV afterglow which radiate about 0.01— 0.1e,/ep)*Frey every decade of time for hours and days after the bursts.", Therefore we expect a bright GeV afterglow which radiate about $0.01$ $0.1 (\epsilon_{e}/\epsilon_B)^{1/2} F_{\rm MeV}$ every decade of time for hours and days after the bursts.393 In this paper when considering EGRET observations. we adopt 200 s after Του. when electrons are in the fast cooling regime. as the duration over which Fc;v is integrated.," In this paper when considering EGRET observations, we adopt 200 s after $T_{90}$, when electrons are in the fast cooling regime, as the duration over which $F_{\rm GeV}$ is integrated."394 analyzed GRBs that were detected by BATSE and observed by EGRET., analyzed GRBs that were detected by BATSE and observed by EGRET.395 Since the field of view of EGRET was much smaller than that of BATSE and the observation was limited by the life time of the spark chamber. EGRET covered only about 100 GRBs out of ~3000 BATSE bursts.," Since the field of view of EGRET was much smaller than that of BATSE and the observation was limited by the life time of the spark chamber, EGRET covered only about 100 GRBs out of $\sim$ 3000 BATSE bursts."396 But this is still a reasonably large number to get statistically meaningful result., But this is still a reasonably large number to get statistically meaningful result.397 The analysis of the prompt burst in EGRET data was performed around the error circles of BATSE bursts for the first 750. and spectral index of —2.4 Is assumed within EGRET window (the upper limits are higher by a factor of z 10 for a spectral index of —1).," The analysis of the prompt burst in EGRET data was performed around the error circles of BATSE bursts for the first $T_{90}$, and spectral index of $-2.4$ is assumed within EGRET window (the upper limits are higher by a factor of $\approx$ 10 for a spectral index of $-1$ )."398 The same analysis was performed for the afterglow phase. for 200 s after Zoo (not including 750).," The same analysis was performed for the afterglow phase, for 200 s after $T_{90}$ (not including $T_{90}$ )."399 measured the fluence of 6 and 12 GRBs. in prompt and afterglow phases respectively.," measured the fluence of 6 and 12 GRBs, in prompt and afterglow phases respectively."400 For all other GRBs only fluence upper limits were obtained in the range 1079-107 erg em., For all other GRBs only fluence upper limits were obtained in the range $10^{-6}$ $10^{-3}$ erg $^{-2}$.401 Here we interpret these results in the framework of the SSC model. which implies that the fluences in BATSE and EGRET bands are likely to be positively correlated through equation (1) CFgarsi=Fyiv and νοκ]=Fo).," Here we interpret these results in the framework of the SSC model, which implies that the fluences in BATSE and EGRET bands are likely to be positively correlated through equation \ref{eq:EGRET402fluence}) ) $F_{\rm BATSE}=F_{\rm MeV}$ and $F_{\rm EGRET}=F_{\rm403GeV}$ )."404 We further assume that the coefficient jj (44 for prompt and jc for afterelow phases) follows some probability distribution function pGp which is independent of Fgayse., We further assume that the coefficient $\eta$ $\eta_{\rm syn}$ for prompt and $\eta_{\rm IC}$ for afterglow phases) follows some probability distribution function $p(\eta)$ which is independent of $F_{\rm BATSE}$.405" We consider a log-normal distribution with the central value µ and standard deviation o: Constraining jj and σ then leads to implications of GRB parameters such as €,. eg. and ~,,. through their relations given in the previous section."," We consider a log-normal distribution with the central value $\mu$ and standard deviation $\sigma$: Constraining $\mu$ and $\sigma$ then leads to implications of GRB parameters such as $\epsilon_{e}$, $\epsilon_B$, and $\gamma_{m}$, through their relations given in the previous section."406 We used the observations to constrain j/ and o by carrying out a maximum likelihood Figure 2. shows the contour plot of the most likely region on the µ-σ plane for Too (top) and 200 s after Zoo data (bottoni) assuming a spectral index of —2.4 (1f the spectral index is —1] then jy increases by +1)., We used the observations to constrain $\mu$ and $\sigma$ by carrying out a maximum likelihood Figure \ref{fig:mu_sig} shows the contour plot of the most likely region on the $\mu$ $\sigma$ plane for $T_{90}$ ) and 200 s after $T_{90}$ data ) assuming a spectral index of $-2.4$ (if the spectral index is $-1$ then $\mu$ increases by $\approx$ 1).407 In that procedure. detection efficiency of EGRET as a function offluence. e(F). is obtained from the distribution of the EGRET upper limits (for undetected GRBs). which ts shown in Figure 3:: re. a cumulative fraction of bursts whose fluence limits are below a given fluence.," In that procedure, detection efficiency of EGRET as a function offluence, $\epsilon (F)$ , is obtained from the distribution of the EGRET upper limits (for undetected GRBs), which is shown in Figure \ref{fig:efficiency}; ; i.e., a cumulative fraction of bursts whose fluence limits are below a given fluence."408 In the case of detected GRBs. on the other hand. the size of the error bars for the fluence is interpreted as measurement accuracy of EGRET.," In the case of detected GRBs, on the other hand, the size of the error bars for the fluence is interpreted as measurement accuracy of EGRET."409 Then. in order to test the consistency of the assumption that equation (11)) fits the data. we carriedout à Monte Carlo simulation that draws," Then, in order to test the consistency of the assumption that equation \ref{eq:eta distribution}) ) fits the data, we carriedout a Monte Carlo simulation that draws"410been recently supported by Bloom et al. (2003)),been recently supported by Bloom et al. \cite{Bloo03}) )411 who based on the optical emission lines suggest that the host is a starburst galaxy. rather than a LINER or a Seyfert 2 galaxy.," who based on the optical emission lines suggest that the host is a starburst galaxy, rather than a LINER or a Seyfert 2 galaxy."412 Bloom et al. (2003)), Bloom et al. \cite{Bloo03}) )413 estimate an unobscured star formation rate (SFR) of 55M... + based on the [OIL] emission line diagnostic method (Kennicutt 1992 ))., estimate an unobscured star formation rate (SFR) of $55 M_{\odot}$ $^{-1}$ based on the ] emission line diagnostic method (Kennicutt \cite{Kenn92}) ).414 Berger et al. (2003)), Berger et al. \cite{Berg03}) )415 and Barnard. et al. (2003)), and Barnard et al. \cite{Barn03}) )416 have recently reported tentative detections of several GRB host galaxies i the sub-millimeter and radio ranges. inferring SFRs of hundreds of Solar masses per year.," have recently reported tentative detections of several GRB host galaxies in the sub-millimeter and radio ranges, inferring SFRs of hundreds of Solar masses per year."417 The most significant detection was from the host galaxy of GRB 000418 for which the sub-mm detection is significant at the 3.607 level and the radio emission from the host again at more than 3o., The most significant detection was from the host galaxy of GRB 000418 for which the sub-mm detection is significant at the $\sigma$ level and the radio emission from the host again at more than $\sigma$.418 Berger et al. (2003)), Berger et al. \cite{Berg03}) )419 also detected another faint radio source only 1.4 aresec from the GRB 000418 host., also detected another faint radio source only 1.4 arcsec from the GRB 000418 host.420 This source is not seen in the optical down to R27., This source is not seen in the optical down to $>$ 27.421 Assuming that the sub-mm and radio emission is due to dust heated by star formation Berger et al. (2003)), Assuming that the sub-mm and radio emission is due to dust heated by star formation Berger et al. \cite{Berg03}) )422 derived SERs of 6904195 and 330075 AL. + respectively. 1.9. much higher than the SFR derived from the optical emission lines by Bloom et al. (2003)).," derived SFRs of $\pm$ 195 and $\pm$ 75 $M_{\odot}$ $^{-1}$ respectively, i.e. much higher than the SFR derived from the optical emission lines by Bloom et al. \cite{Bloo03}) )."423 In this paper we present a comprehensive multicolour study of the GRB 000418 host galaxy. similar to those performed for the GRB 000210 (Gorosabel et al. 2003))," In this paper we present a comprehensive multicolour study of the GRB 000418 host galaxy, similar to those performed for the GRB 000210 (Gorosabel et al. \cite{Goro03}) )"424 and GRB 990712 (Christensen et al. 2003)), and GRB 990712 (Christensen et al. \cite{Chri03}) )425 host galaxies., host galaxies.426 The aim of the analysis is to determine the properties of the stellar populations dominating the optical/near-IR light from the host galaxy and the amount of extinction due to dust in the interstellar medium (ISM) of thehost., The aim of the analysis is to determine the properties of the stellar populations dominating the optical/near-IR light from the host galaxy and the amount of extinction due to dust in the interstellar medium (ISM) of thehost.427 Other multi-colour host galaxies studies to date (Sokolov et al. 2001..," Other multi-colour host galaxies studies to date (Sokolov et al. \cite{Soko01},"428 Chary et al. 2002.. ," Chary et al. \cite{Char02}, ,"429Gorosabel et al. 2001) , Gorosabel et al. \cite{Goro01}) )430have been limited to a smaller number of bands., have been limited to a smaller number of bands.431" Throughout. the assumed cosmology will be O4=0.7. Oy,—0.3 and Jy—G5 kms | |."," Throughout, the assumed cosmology will be $\Omega_{\Lambda} = 0.7$, $\Omega_{M} = 0.3$ and $H_0= 65$ km $^{-1}$ $^{-1}$."432 At the spectroscopic redshift of the host galaxy ἐς= 1.118). the look back time is 8.78 Gyr (z of the present age) and the luminosity distance is 8.17 Gpe.," At the spectroscopic redshift of the host galaxy $z=1.118$ ), the look back time is 8.78 Gyr $\approx$ of the present age) and the luminosity distance is 8.17 Gpc."433 The physical transverse size of one aresec at +=1.115 corresponds to 8.83 kpe., The physical transverse size of one arcsec at $z=1.118$ corresponds to 8.83 kpc.434 We have used ground-based and space optical/NIR resources in order to sample the host galaxy SED in as many filters as possible over a wide spectral range (see Table 1))., We have used ground-based and space optical/NIR resources in order to sample the host galaxy SED in as many filters as possible over a wide spectral range (see Table \ref{table1}) ).435 The observations were performed 299-676 days after the gamma-ray event., The observations were performed 299–676 days after the gamma-ray event.436 At these times the contribution of the afterglow is negligible and does not have any significant impact on the derived host galaxy SED., At these times the contribution of the afterglow is negligible and does not have any significant impact on the derived host galaxy SED.437 Given the compactness of the GRB 000418 host galaxy. its ground-based NIR and optical profiles are consistent with that of field stars.," Given the compactness of the GRB 000418 host galaxy, its ground-based NIR and optical profiles are consistent with that of field stars."438 This fact has been checked by comparing the growth curve (from 0.5 to 4 times their full width half maximum: FWHM hereafter) of the host and the secondary stars used for photometric calibration., This fact has been checked by comparing the growth curve (from 0.5 to 4 times their full width half maximum; FWHM hereafter) of the host and the secondary stars used for photometric calibration.439 Thus. for the ground-based optical and NIR pixel scales used in the present study ϱ071IS/pix. achieved with ISAAC). the GRB 000418 host galaxy can safely be assumed to be a point-source.," Thus, for the ground-based optical and NIR pixel scales used in the present study $\ge 0\farcs148$ /pix, achieved with ISAAC), the GRB 000418 host galaxy can safely be assumed to be a point-source."440 Therefore. considering that the relative photometry ts independent of the aperture radius. the (BVRIZ.N.-band magnitudes shown in Table 2 are based on circular aperture (PHOT running under IRAF-)) with no aperture corrections.," Therefore, considering that the relative photometry is independent of the aperture radius, the $UBVRIZJ_sK_s$ -band magnitudes shown in Table \ref{table2} are based on circular aperture (PHOT running under ) with no aperture corrections."441 DVRIZ-band frames were taken with ALFOSC at the 2.5- Nordic Optical Telescope (2.5NOT)., $BVRIZ$ -band frames were taken with ALFOSC at the 2.5-m Nordic Optical Telescope (2.5NOT).442 The ALFOSC detector is à «2045 Thinned Loral CCD providing a pixel scale of Y1589/pix., The ALFOSC detector is a $\times$ 2048 Thinned Loral CCD providing a pixel scale of $0\farcs189$ /pix.443 C-band observations were carried out with the 3.6-n ESO telescope (3.6ESO) equipped with EFOSC2. covering a field of view (FOV) of 5/5«5/5.," $U$ -band observations were carried out with the 3.6-m ESO telescope (3.6ESO) equipped with EFOSC2, covering a field of view (FOV) of $5\farcm5 \times4445\farcm5$."445 These observations were carried out in 242 binning mode. providing a pixel scale of Y/31 Upix.," These observations were carried out in $\times$ 2 binning mode, providing a pixel scale of $0\farcs314$ /pix."446 The optical data were reduced in à standard manner (overscan. bias subtraction. and division by a normalised flat field).," The optical data were reduced in a standard manner (overscan, bias subtraction, and division by a normalised flat field)."447 The (BV Ri-band calibration was based on the calibration given by Henden (2000))., The $UBVRI$ -band calibration was based on the calibration given by Henden \cite{Hend00}) ).448 The Z-band calibration was carried out observing the spectro-photometric standard star Feige 66 (Oke 1990)) with the 2.5NOTC-ALFOSC) at an airmass almost identical (Asec(:)= 0.03) to the GRB field.," The $Z$ -band calibration was carried out observing the spectro-photometric standard star Feige 66 (Oke \cite{Oke90}) ) with the 2.5NOT(+ALFOSC) at an airmass almost identical $\Delta \sec(z) =4490.03$ ) to the GRB field."450 Fig., Fig.451 | shows a deep V -band image of the GRB field taken with the 2.5NOT(+ALFOSC)., \ref{2.5NOT} shows a deep $V$ -band image of the GRB field taken with the 2.5NOT(+ALFOSC).452" The NIR (7, and A,-band) observations were acquired with the UTI of the 8.2-m VLT (8.2VLT) equipped with ISAAC. allowing us to cover a 2/5« FOV with a pixel scale of 071 Ls/pix."," The NIR $J_s$ and $K_s$ -band) observations were acquired with the UT1 of the 8.2-m VLT (8.2VLT) equipped with ISAAC, allowing us to cover a $2\farcm4535 \times 2\farcm 5$ FOV with a pixel scale of $0\farcs148$ /pix."454 In Table | we provide the observing log of our optical and NIR observations., In Table \ref{table1} we provide the observing log of our optical and NIR observations.455" The calibration was based on observations of the standard stars S301-D (7, band) and S860-D (NV. band: Persson et al. 1998)).", The calibration was based on observations of the standard stars S301-D $J_s$ band) and S860-D $K_s$ band; Persson et al. \cite{Pers98}) ).456 Due to the lack of J.-band calibration data for S301-D we assumed J)=J.., Due to the lack of $J_s$ -band calibration data for S301-D we assumed $J=J_s$.457" Extinction coefficients of 0.06 and 0.07 mag per airmass (ESO Paranal) where assumed for ./, and A. respectively."," Extinction coefficients of 0.06 and 0.07 mag per airmass (ESO Paranal) where assumed for $J_s$ and $K_s$, respectively."458" Given that the airmass difference between the GRB field and the standard was Asec(:)=0. Lin both -/, and A. the introduced airmass correction is well below our measurementerror on the magnitude of the host galaxy."," Given that the airmass difference between the GRB field and the standard was $\Delta \sec(z) = 0.4$ in both $J_s$ and $K_s$, the introduced airmass correction is well below our measurementerror on the magnitude of the host galaxy."459 The photometric calibration has been tested performing photometry of the 2MASS star at RA. DEC (J2000) = 12:25:14.47. 20:05:49.8 which is present on the ISAAC images taken for the host.," The photometric calibration has been tested performing photometry of the 2MASS star at RA, DEC (J2000) = 12:25:14.47, 20:05:49.8 which is present on the ISAAC images taken for the host."460" Using the S301-D star we derived for this source Jo=15.51+0.01 and A,=LL8T+ 0.01.", Using the S301-D star we derived for this source $J_s = 15.51 \pm 0.01$ and $K_s= 14.87 \pm 0.01$ .461 Both magnitudes agree with the ones given in the 2MASS Second, Both magnitudes agree with the ones given in the 2MASS Second462of macroscopic galactic parameters.,of macroscopic galactic parameters.463 This is confirmed by our results., This is confirmed by our results.464 The huge dispersion in c-drop sizes can be related to a reduction of their size when they become old or to the different extents of the cool gas volume at the moment it reached the critical density., The huge dispersion in $\sigma$ -drop sizes can be related to a reduction of their size when they become old or to the different extents of the cool gas volume at the moment it reached the critical density.465 As we have shown in Section 4.2. o-drop hosts have in more than of the cases a clear inner dust system of spiral arms that probably trace the path of the inflowing material.," As we have shown in Section 4.2, $\sigma$ -drop hosts have in more than of the cases a clear inner dust system of spiral arms that probably trace the path of the inflowing material."466 In contrast. less than of the control sample galaxies show such a clear spiral pattern.," In contrast, less than of the control sample galaxies show such a clear spiral pattern."467 As there are no significant differences in bar characteristics between the two samples. it is safe to say that there is some correlation between c-drops and inner dust spiral arms.," As there are no significant differences in bar characteristics between the two samples, it is safe to say that there is some correlation between $\sigma$ -drops and inner dust spiral arms."468" This dust could trace the inflowing gas fuelling the o-drop. or it could be the product of relatively recent star-formation,"," This dust could trace the inflowing gas fuelling the $\sigma$ -drop, or it could be the product of relatively recent star-formation."469 We have also found that o-drop galaxies often (in of the cases) have a nuclear star-forming ring visible in with a size between ppc and kkpe and that there are no star-forming rings at those scales in the control sample., We have also found that $\sigma$ -drop galaxies often (in of the cases) have a nuclear star-forming ring visible in with a size between pc and kpc and that there are no star-forming rings at those scales in the control sample.470 It thus seems reasonable to say that the rings are a manifestation of the same phenomena causing the o-drops., It thus seems reasonable to say that the rings are a manifestation of the same phenomena causing the $\sigma$ -drops.471 The stars that compose the o-drops can be formed in such nuclear starburst rings. usually due to shock focusing of gas near the location of one or more inner Lindblad resonances (ILRs: Schwarz 1984; Athanassoula 1992: Knapen et al.," The stars that compose the $\sigma$ -drops can be formed in such nuclear starburst rings, usually due to shock focusing of gas near the location of one or more inner Lindblad resonances (ILRs; Schwarz 1984; Athanassoula 1992; Knapen et al."472" 1995. Heller Shlosman 1996: see also Shlosman 1999, Knapen 2005 for reviews)."," 1995, Heller Shlosman 1996; see also Shlosman 1999, Knapen 2005 for reviews)."473 The ILRs are thought to be related to bars or interactions between close galaxies., The ILRs are thought to be related to bars or interactions between close galaxies.474 Four of our five galaxies with star-forming rings are barred., Four of our five galaxies with star-forming rings are barred.475 The fifth. NGC 6503 is so edge-on that a bar may be difficult to detect.," The fifth, NGC 6503 is so edge-on that a bar may be difficult to detect."476 We have found that o-drop galaxies are more often Seyfert than LINER hosts., We have found that $\sigma$ -drop galaxies are more often Seyfert than LINER hosts.477 On the other hand. control galaxies more often have LINER emission than Seyfert emission.," On the other hand, control galaxies more often have LINER emission than Seyfert emission."478 Seyfert galaxies are more energetic phenomena. which probably implies a more continuous or more recent inflow of material to maintain the central activity.," Seyfert galaxies are more energetic phenomena, which probably implies a more continuous or more recent inflow of material to maintain the central activity."479 The dust spiral arms that are frequent in c-drop galaxies are a plausible method to feed the central black hole., The dust spiral arms that are frequent in $\sigma$ -drop galaxies are a plausible method to feed the central black hole.480 It 1s therefore possible to postulate that both o-drops and Seyferts are due to a more efficient feeding of the inner parts of galaxies., It is therefore possible to postulate that both $\sigma$ -drops and Seyferts are due to a more efficient feeding of the inner parts of galaxies.481 As Seyferts are shorter lived than σ- (few 10° yyrs versus GGyr) the correlation ean only hold partially., As Seyferts are shorter lived than $\sigma$ -drops (few $10^{8}$ yrs versus Gyr) the correlation can only hold partially.482 We thus suggest the following model for the creation of à c-drop., We thus suggest the following model for the creation of a $\sigma$ -drop.483 Gas is driven inwards by spiral arms and maybe by bars., Gas is driven inwards by spiral arms and maybe by bars.484 The gas is focused by an ILR into a dynamically cold ring where the density increases until it reaches a critical value when star formation starts., The gas is focused by an ILR into a dynamically cold ring where the density increases until it reaches a critical value when star formation starts.485 These stars are dynamically cold and start to create a c-drop., These stars are dynamically cold and start to create a $\sigma$ -drop.486 Alternatively. if there is no ILR. gas can be focused in a cold nuclear dise that also starts to create stars when it reaches the critical density.," Alternatively, if there is no ILR, gas can be focused in a cold nuclear disc that also starts to create stars when it reaches the critical density."487 In both cases friction causes some gas to lose enough angular momentum to drift inwards to feed the super-massive black hole., In both cases friction causes some gas to lose enough angular momentum to drift inwards to feed the super-massive black hole.488 This gas can be traced with the dust related to the relatively recent star formation., This gas can be traced with the dust related to the relatively recent star formation.489 It is entirely possible that c-drops are exclusively created in rings that disappear due to the transient nature of the bars and/or the inflow., It is entirely possible that $\sigma$ -drops are exclusively created in rings that disappear due to the transient nature of the bars and/or the inflow.490 As a c-drop is a long-lived feature (as shown by modelling by Wozniak&Champavert2006.. and as indicated by the high fraction of galaxies that host them) it would survive long after the ring has disappeared.," As a $\sigma$ -drop is a long-lived feature (as shown by modelling by \cite{WO06}, and as indicated by the high fraction of galaxies that host them) it would survive long after the ring has disappeared."491 These conclusions should be strengthened in future work by improving the statistics and by further numerical modelling., These conclusions should be strengthened in future work by improving the statistics and by further numerical modelling.492the instantaneous SF model reaches its maximum value very quickly with the evolution of the thin disc.,the instantaneous SF model reaches its maximum value very quickly with the evolution of the thin disc.493 The maximum value is significantly higher than that in another three SF models since the initial SF rate in the instantaneous SF model is more than ~30 times that in other SF models., The maximum value is significantly higher than that in another three SF models since the initial SF rate in the instantaneous SF model is more than $\sim$ 30 times that in other SF models.494 The enhanced SF model produces the highest present rate of SNIa in the four SF models., The enhanced SF model produces the highest present rate of SNIa in the four SF models.495 This is because the present SF contributes the majority of mergers., This is because the present SF contributes the majority of mergers.496" Again, from reffig.nia, , ο ο ο ο SN Toke have mittentwa ihepjsodes"," Again, from \\ref{fig_snia}, continuous SF is also the reason why the rate of SNIa keeps growing with the the evolution of the thin disc."497" ruleο... sed” daySN Iarateisabout2.19, 5.07, and4.34x107 quasi-exponentialyr! for the instantaneous, constant, enhanced,4.92, and SF models respectively."," The present-day SNIa rate is about 2.19, 4.92, 5.07, and 4.34 $\times10^{-4}$ $^{-1}$ for the instantaneous, constant, enhanced, and quasi-exponential SF models respectively."498 We have to emphasize that we here only take into account the merger channel for the rate of thin-disc SNIa., We have to emphasize that we here only take into account the merger channel for the rate of thin-disc SNIa.499" This leads to our results being less than the occurrence rate inferred from observations, ~4x1073 1 (2),, as a few other formation channels would contribute to the observed occurrence rate (???).."," This leads to our results being less than the occurrence rate inferred from observations, $\sim4\times10^{-3}$ $^{-1}$ \citep{Cappellaro97}, as a few other formation channels would contribute to the observed occurrence rate \citep{Yungelson00,Han04,Hachisu08}."500 We note that the occurrence (or birth) rate of SNIa at fais. Gyr in our simulations is smaller than in some other recent theoretical studies ??)) by a factor of 1.5—4.," We note that the occurrence (or birth) rate of SNIa at $t_{\rm501 disc}=10$ Gyr in our simulations is smaller than in some other recent theoretical studies \citet{Ruiter09,Mennekens10}) ) by a factor of $1.5-4$."502" Since all these studies (including ours) adopt the same IMF for the primary main sequence stars, the same flat distribution for initial mass ratio, and the same initial eccentricity distribution, we infer that the differences between our study and others are mainly caused by the initial distribution of orbital separations and by the treatment of mass transfer and mass loss during binary evolution."," Since all these studies (including ours) adopt the same IMF for the primary main sequence stars, the same flat distribution for initial mass ratio, and the same initial eccentricity distribution, we infer that the differences between our study and others are mainly caused by the initial distribution of orbital separations and by the treatment of mass transfer and mass loss during binary evolution."503" Both. ? and? adopt a distribution of the orbital separation anoa!(?) when a<10° Ro, while we use a distribution of xa9? (??) fora <10 Ro, and cca! for 10<ας5.75x108 Ro."," Both \citet{Ruiter09} and \citet{Mennekens10} adopt a distribution of the orbital separation $\frac{{\rm d} n}{{\rm d} a}\propto a^{-1}$ \citep{abt83} when $a\leqslant10^{5}$ $R_{\odot}$, while we use a distribution of $\propto a^{0.2}$ \citep{Griffin85,Han98} for $a\leqslant10$ $R_{\odot}$, and $\propto a^{-1}$ for $10\leqslant a\leqslant5.75\times10^{6}$ $R_{\odot}$."504 This may lead to the generation of more close DDs in their models than our model., This may lead to the generation of more close DDs in their models than our model.505" For common-envelope ejection, we adopt the -algorithm (?) and fix ;=1.5."," For common-envelope ejection, we adopt the $\gamma$ -algorithm \citep{Nelemans05} and fix $\gamma=1.5$."506 Other studies use the a-algorithm (?).., Other studies use the $\alpha$ -algorithm \citep{Webbink84}.507" Increasing the common-envelope ejection efficiency (for example, inreasing aX=0.5—1.0) results in more binaries going through a double common-envelope phase to become close DDs, and hence forming potential SNIa progenitors."," Increasing the common-envelope ejection efficiency (for example, inreasing $\alpha\lambda=0.5\rightarrow1.0$ ) results in more binaries going through a double common-envelope phase to become close DDs, and hence forming potential SNIa progenitors."508 A close comparison between the studies implies that the common-envelope ejection efficiency is lower in our calculation than in others., A close comparison between the studies implies that the common-envelope ejection efficiency is lower in our calculation than in others.509" In the most extreme case, the ? SNIa birth rate from the DD merger channel is higher than ours by a factor of ~10 because they adopted a very high common-envelope ejection efficiency (αλ=2.0) and a higher SF rate (~8 Moyr7')."," In the most extreme case, the \citet{Yungelson10} SNIa birth rate from the DD merger channel is higher than ours by a factor of $\sim$ 10 because they adopted a very high common-envelope ejection efficiency $\alpha\lambda=2.0$ ) and a higher SF rate $\sim$ 8 $M_{\odot}$ $^{-1}$ )."510 A significant difference between the studies concerns the star-formation history for early-type (elliptical) galaxies., A significant difference between the studies concerns the star-formation history for early-type (elliptical) galaxies.511 Both ? and ? approximate “instantaneous star formation” by a 6-function star-burst., Both \citet{Ruiter09} and \citet{Mennekens10} approximate “instantaneous star formation” by a $\delta$ -function star-burst.512" Our approximation of “instantaneous SF"" assumes that the initial star-burst lasted a few million years at a constant rate.", Our approximation of “instantaneous SF” assumes that the initial star-burst lasted a few million years at a constant rate.513 This leads to the peak value of the SNIa rate occurring rather later in our models than in the studies in ? and ?.., This leads to the peak value of the SNIa rate occurring rather later in our models than in the studies in \citet{Ruiter09} and \citet{Mennekens10}.514" However, none of these approximations may reflect the true SF history of early- galaxies; recent studies indicate that bright early-type galaxies show signs of current star formation (?) and, by inferrence, may αι.experience thet}te luteos"," However, none of these approximations may reflect the true SF history of early-type galaxies; recent studies indicate that bright early-type galaxies show signs of current star formation \citep{Yi05} and, by inferrence, may have experienced intermittent star-burst episodes."515tar-bu: ase 6 SES a, The case of ccSN is easier than that of SNIa.516" Their progenitors, main sequence stars with mass greater than 8 Mo, evolve very quickly on a timescale «20 Myr."," Their progenitors, main sequence stars with mass approximately greater than 8 $M_{\odot}$, evolve very quickly on a timescale $\lesssim$ 20 Myr."517" This is even shorter than the tg of ONeMg+X (see soallthenew bornccS Nresults rogenitor)), fromS FafterafewtensM yr(seecriterionin& secys))."," This is even shorter than the $t_{\rm MS}$ of ONeMg+X (see \\ref{fig_progenitor}) ), so all the new-born ccSN results from SF after a few tens Myr (see criterion in \\ref{sec_ps}) )."518ThentheS FratecanberegardedasaconstantS(taisc)., Then the SF rate can be regarded as a constant $S(t_{\rm disc}$ ).519" From refeq,irthrate, , wehavethebirthrateofthesupernovaeSNR(taisc) . S(taisc), where Necgn is the number of stars going supernova in a calculation of a sample of main sequence stars, Mtotal represents the total mass of the sample, and taisc=10 Gyr."," From \\ref{eq_birthrate}, we have the birth rate of the supernovae $SNR(t_{\rm disc})\approx[N_{\rm ccSN}/m_{\rm total}]\cdot S(t_{\rm disc})$ , where $N_{\rm ccSN}$ is the number of stars going supernova in a calculation of a sample of main sequence stars, $m_{\rm total}$ represents the total mass of the sample, and $t_{\rm disc}=10$ Gyr."520" Using the same parameters as in the simulation in this paper, we obtain Necsn/Mtotal&24.5X107? and so SNR(10)z 2.34, 2.7 and 1.58 century! for the three Mg!continuous SF models."," Using the same parameters as in the simulation in this paper, we obtain $N_{\rm ccSN}/m_{\rm total}\approx4.5\times10^{-3}$ $M_{\odot}^{-1}$ and so $SNR(10)\approx$ 2.34, 2.7 and 1.58 $^{-1}$ for the three continuous SF models."521" These values are consistent with the observations of 1.2 —3 century! (2222?) in terms of the mass of ?9 AI, the number of massive stars in the localized HII region and the extragalactic supernova scaled calculation."," These values are consistent with the observations of 1.2 $-$ 3 $^{-1}$ \citep{Smith78,vandenbergh94,McKee97,Timmes97,Diehl06} in terms of the mass of $^{26}$ Al, the number of massive stars in the localized HII region and the extragalactic supernova scaled calculation."522 Our simulation is based only on theory with some simple assumptions., Our simulation is based only on theory with some simple assumptions.523" The final outputs, the rates and numbers, are derived from two key inputs which are the SF rate and the contribution functions for the evolved stars."," The final outputs, the rates and numbers, are derived from two key inputs which are the SF rate and the contribution functions for the evolved stars."524" On the other hand, if we know the rates and numbers of objects in a galaxy, especially some types of exotic stars, we are able to do a deconvolution by setting the Observed rate as an input parameter to derive the SF rate at specific"," On the other hand, if we know the rates and numbers of objects in a galaxy, especially some types of exotic stars, we are able to do a deconvolution by setting the observed rate as an input parameter to derive the SF rate at specific"525 (Fe/Il]X2.5) Wallersteietal.(1963) (Soubira," $\feh\lesssim-2.5$ $V=6.2$ \cite{wallerstein63} \cite[e.g.,][]{sneden73:hd122563,lambert74,ryan96,cayrel04}. \citep{soubiran10},"526uetal.2010).. logye=1.25£0.22. 2008)..," $\teff=4602\pm72$ $\logg=1.25\pm0.22$ $\feh=-2.64\pm0.12$ \citep[e.g.,][]{nesis92,rutten04} \citep[e.g.,][]{allende02,gray05,dravins08,kdwarfs-p1}."527 features so that the line cores of disk-iutegrated profiles are shifted with respect to their rest waveleneths aud the line profiles are asvunuctric (e.g..Draviusctal.1981).," features so that the line cores of disk-integrated profiles are shifted with respect to their rest wavelengths and the line profiles are asymmetric \citep[e.g.,][]{dravins81}."528. The latter are granulation signatures that can be observed both in the solar spectra aud in the spectra of distant stars., The latter are granulation signatures that can be observed both in the solar spectrum and in the spectra of distant stars.529 νάνοςναας simulations provide the most realistic approach to the modeliug of surface inhomogencitics (e.g...Nordluud1982:Steffen&Frevtag1991:Lud-wieetal.2009:Nordlund 2009).. although lydrodvuamic modeling is also possible (e.g...Dravius 1990).," Hydrodynamic simulations provide the most realistic approach to the modeling of surface inhomogeneities \citep[e.g.,][]{nordlund82,steffen91,ludwig09,nordlund09:sun}, although non-hydrodynamic modeling is also possible \citep[e.g.,][]{dravins90}."530. Threc-dinensional model atmiospheres for the Sun. Procvon. aud IE-dwarfs have beeu tested agaist high quality spectroscopic data (Asplundetal.2000:Al-leudePrietoetal.2002:Ramirez 20093.. showing in general good agreeimeut.," Three-dimensional model atmospheres for the Sun, Procyon, and K-dwarfs have been tested against high quality spectroscopic data \citep{asplund00:iron_shapes,allende02,kdwarfs-p2}, showing in general good agreement."531" These tests are crucial because they validate the models and allow us to use them with coufidence to estimate the impact of ""3D effects onu stellar abundance anc parameter determiunatious (Asplund2005).", These tests are crucial because they validate the models and allow us to use them with confidence to estimate the impact of “3D effects” on stellar abundance and parameter determinations \citep{asplund05:review}.532 The most extreme abuudance correctious proposed in connection with eranulation effects occur at low netallicities., The most extreme abundance corrections proposed in connection with granulation effects occur at low metallicities.533 This is iainly because the difference between the 1D aud 3D predictious for thle mean temperature of the outer lavers of mctal-poor stars ds very laree. on the order of IKI&. (6.5...Colletetal.2006.2007:ConzálezUerndudezetal. 2010)..," This is mainly because the difference between the 1D and 3D predictions for the mean temperature of the outer layers of metal-poor stars is very large, on the order of K \citep[e.g.,][]{collet06,collet07,gonzalez-hernandez10:3d}."534 This laree temperature difference is very important for the formation of features that are temperature sensitive., This large temperature difference is very important for the formation of features that are temperature sensitive.535 For example. Colletetal.(2009). report 3D corrections of ~ lddex (.6.. one order of magnitude)/ for the abundance of nitrogeu aud oxveen as deteriuned frou molecular features. simular to the corrections required for the iron abundance determined from low excitation potecutial Hines.," For example, \cite{collet09} report 3D corrections of $\sim-1$ dex (i.e., one order of magnitude) for the abundance of nitrogen and oxygen as determined from molecular features, similar to the corrections required for the iron abundance determined from low excitation potential lines."536 They also estimate a correction of about dclex, They also estimate a correction of about dex537 They also estimate a correction of about dclex:, They also estimate a correction of about dex538"this that our orbit fitting code simultaneously fits a simple spin-up model and an orbital radial velocity model where P,»; is the detected spin period of the pulsar, c is the speed of light and ν( is the binary radial velocity.","this that our orbit fitting code simultaneously fits a simple spin-up model and an orbital radial velocity model where $P_{obs}$ is the detected spin period of the pulsar, $c$ is the speed of light and $v(t)$ is the binary radial velocity."539 v(t) is calculated using the IDL routine from the which uses the procedure of Hilditch(2001) and employs the method of Mikkola(1987) to solve Kepler’s equation., $v(t)$ is calculated using the IDL routine from the which uses the procedure of \citet{2001icbs.book.....H} and employs the method of \citet{1987CeMec..40..329M} to solve Kepler's equation.540 The spin of the pulsar is given by P(t) which incorporates the spin-up component and is defined below: where P and P are the spin-up and change in spin-up of the NS., The spin of the pulsar is given by $P(t)$ which incorporates the spin-up component and is defined below: where $\dot{P}$ and $\ddot{P}$ are the spin-up and change in spin-up of the NS.541 We perform a least-squares fit on the data and return the best fitting parameters., We perform a least-squares fit on the data and return the best fitting parameters.542 The full spin-up component shown in equation 4 was fit to every outburst., The full spin-up component shown in equation \ref{equ:spin} was fit to every outburst.543" However, no data set allowed for sufficient statistics in the fit to justify a P component and hence it was removed."," However, no data set allowed for sufficient statistics in the fit to justify a $\ddot{P}$ component and hence it was removed."544 The results of our fits are presented in the next section., The results of our fits are presented in the next section.545" In order to obtain a meaningful fit, the period measurements used in the fitting procedure were restricted to detections greater than a particular significance."," In order to obtain a meaningful fit, the period measurements used in the fitting procedure were restricted to detections greater than a particular significance."546" This threshold was not kept constant for every source because the luminosity and duration of each outburst was different, meaning fits improved or worsened by including data above different thresholds."," This threshold was not kept constant for every source because the luminosity and duration of each outburst was different, meaning fits improved or worsened by including data above different thresholds."547" As such, a significance threshold was chosen to minimise the (16) errors in the binary parameters for each fit."," As such, a significance threshold was chosen to minimise the $\sigma$ ) errors in the binary parameters for each fit."548 The chosen value is stated in the relevant section., The chosen value is stated in the relevant section.549 SXP6.85 was first detected in 2003 by as a ss pulsed X-ray source (Corbetetal. 2003).., SXP6.85 was first detected in 2003 by as a s pulsed X-ray source \citep{cor03}. .550" It was later detected in a 2006 observation at the position R.A. = 0102""53:1, dec."," It was later detected in a 2006 observation at the position R.A. = $01^{h}02^{m}53\fs1$ , dec."551 = —72°44'33/0 (J2000.0)., = $-72^{\circ}44^{'}33\farcs0$ (J2000.0).552" This detection led to the identification of a V=14.6 optical counterpart (Haberletal.2008),, allowing it to be classified as an HMXB."," This detection led to the identification of a V=14.6 optical counterpart \citep{hab08}, allowing it to be classified as an HMXB."553 Follow-up work by McBrideetal.(2008) classified the counterpart as an VV emission line star., Follow-up work by \citet{mcbride08} classified the counterpart as an V emission line star.554" In subsequent years it has been detected on 5 distinct occasions, coinciding with times when the counterpart was optically bright (Townsendetal.2010)."," In subsequent years it has been detected on 5 distinct occasions, coinciding with times when the counterpart was optically bright \citep{town10}."555. McGowanetal. find that the optical flux varies by ~00.5 magnitudes with a period of 620+18 dd. They associate this with the growth and decay of the circumstellar disk., \citet{kem08} find that the optical flux varies by $\sim$ 0.5 magnitudes with a period of $620 \pm 18$ d. They associate this with the growth and decay of the circumstellar disk.556" Those authors also show that the source gets redder as it gets brighter, suggestive of a low inclination system, and propose a low eccentricity based on comparison to other systems."," Those authors also show that the source gets redder as it gets brighter, suggestive of a low inclination system, and propose a low eccentricity based on comparison to other systems."557" Until now, the orbital period of this system was not known for certain."," Until now, the orbital period of this system was not known for certain."558" Analysis of optical light curves by McGowanetal.(2008) and Schmidtke&Cowley(2007) showed hints of periodicities at 114.1+0.6 dd and 24.8+0.1 dd, but neither could be confirmed as the orbital period of the system."," Analysis of optical light curves by \citet{kem08} and \citet{schmit07} showed hints of periodicities at $114.1 \pm 0.6$ d and $24.8 \pm 0.1$ d, but neither could be confirmed as the orbital period of the system."559" The dd period was closer to the expected orbital period based on the Corbet diagram (Corbet, 1986)., but the dd period was reinforced by the detection of a dd period in the X-ray light curve (Galacheetal.2008)."," The d period was closer to the expected orbital period based on the Corbet diagram \citep{cor86}, but the d period was reinforced by the detection of a d period in the X-ray light curve \citep{gal08}."560". To try and resolve this issue, the two longest Type II X-ray outbursts were fit with our orbital model."," To try and resolve this issue, the two longest Type II X-ray outbursts were fit with our orbital model."561" The outburst occuring around MJD 54800 turned out to be too sparsely covered, meaning we were unable to get an acceptable fit."," The outburst occuring around MJD 54800 turned out to be too sparsely covered, meaning we were unable to get an acceptable fit."562" However, the outburst beginning on MJD 55435 (Table 1)) was sampled very well thanks to dedicated pointings at the position of the source in addition to our regular monitoring."," However, the outburst beginning on MJD 55435 (Table \ref{tab:outbursts}) ) was sampled very well thanks to dedicated pointings at the position of the source in addition to our regular monitoring."563 The period evolution is shown in Fig., The period evolution is shown in Fig.564 1 with the best model fit overplotted., \ref{fig:orbit2} with the best model fit overplotted.565 The data used in the fit were cut at the %(4 σ) significance level., The data used in the fit were cut at the $\%$ $\sigma$ ) significance level.566 The very clear changes in the spin period allowed the radial velocity of the NS to be found quite simply., The very clear changes in the spin period allowed the radial velocity of the NS to be found quite simply.567" Unfortunately, we were unable to fit both outbursts simultaneously as the amount of spin-down happening between the outbursts is unknown."," Unfortunately, we were unable to fit both outbursts simultaneously as the amount of spin-down happening between the outbursts is unknown."568 The solution is presented in Table 2 and is the best fit we have in our sample., The solution is presented in Table \ref{tab:orbit2} and is the best fit we have in our sample.569" We propose 21.9+0.1 dd to be the true orbital period of this system, placing it nicely in the BeXRB region of the Corbet diagram."," We propose $21.9\pm0.1$ d to be the true orbital period of this system, placing it nicely in the BeXRB region of the Corbet diagram."570 The other orbital parameters andthe nature of the dd and dd periods are discussed in section 4., The other orbital parameters andthe nature of the d and d periods are discussed in section 4.571 SXP2.37 wasdiscovered by SAS 3 observations in 1977 as a highly variable X-ray source., SXP2.37 wasdiscovered by SAS 3 observations in 1977 \citep{li77} as a highly variable X-ray source.572"than those of of the BAT-triggered bursts, as shown in Fig. Hl.","than those of of the BAT-triggered bursts, as shown in Fig. \ref{fig:XrayOpticalPercentile}."573 The X-ray afterglow of GRB 090510A is brighter than (69%)) of the BAT extended emission (short) bursts., The X-ray afterglow of GRB 090510A is brighter than ) of the BAT extended emission (short) bursts.574 We compared the optical flux in counts s! at 70 ks to 103 bursts with UVOT afterglows included in The Second Telescope GRB Afterglow Catalog Ultra-Violet/Optical(P. Roming et al., We compared the optical flux in counts $^{-1}$ at 70 ks to 103 bursts with UVOT afterglows included in The Second Ultra-Violet/Optical Telescope GRB Afterglow Catalog (P. Roming et al.575" 2010, in preparation)."," 2010, in preparation)."576 All light curves were normalized to the filter and extrapolated to 70 ks (if for our comparison., All light curves were normalized to the filter and extrapolated to 70 ks (if necessary) for our comparison.577" Our preliminary results,necessary) shown in Fig."," Our preliminary results, shown in Fig. \ref{fig:XrayOpticalPercentile},"578" indicate that the optical afterglows of long LAT Hl,bursts are brighter than of BAT-triggered optical afterglows, with GRB 0909264 falling in the top of optical afterglow brightness."," indicate that the optical afterglows of long LAT bursts are brighter than of BAT-triggered optical afterglows, with GRB 090926A falling in the top of optical afterglow brightness."579" Additionally, GRB 090510A is one of only two extended emission GRBs, or one of five short GRBs, still detected by the UVOT at 70 ks."," Additionally, GRB 090510A is one of only two extended emission GRBs, or one of five short GRBs, still detected by the UVOT at 70 ks."580" Regardless of which category (short or extended emission) GRB 090510A belongs to, it is brighter than ~90% of other short/extended emission optical afterglows."," Regardless of which category (short or extended emission) GRB 090510A belongs to, it is brighter than $\sim$ of other short/extended emission optical afterglows."581" We have presented the and observations of GRB 0909264, a recent LAT-detected GRB with a bright, long lived optical afterglow detected by UVOT."," We have presented the and observations of GRB 090926A, a recent LAT-detected GRB with a bright, long lived optical afterglow detected by UVOT."582" We have compared this burst, to other LAT-detected and BAT bursts in an attempt to show whether the GRBs detected by the LAT are simply brighter than the average BAT-triggered GRB or whether they represent a new type of GRB that commonly exhibit bright, long duration optical afterglows due to some form of energy injection."," We have compared this burst, to other LAT-detected and BAT bursts in an attempt to show whether the GRBs detected by the LAT are simply brighter than the average BAT-triggered GRB or whether they represent a new type of GRB that commonly exhibit bright, long duration optical afterglows due to some form of energy injection."583 We find that the LAT-detected bursts are generally brighter than their BAT-triggered counterparts., We find that the LAT-detected bursts are generally brighter than their BAT-triggered counterparts.584" We find that their fluence is consistently higher than the ‘average’ BAT burst, and that their X-ray and UV/optical afterglows are brighter than ~80% of BAT GRBs."," We find that their fluence is consistently higher than the `average' BAT burst, and that their X-ray and UV/optical afterglows are brighter than $\sim$ of BAT GRBs."585" Although we are working with a small sample of LAT bursts, and therefore suffer the consequences of small number statistics, our preliminary results indicate that LAT bursts exhibit bright late time X-ray and UV/optical afterglows because they are brighter at all wavelengths than the ‘average’ burst, assuming the higher than average fluence can be extrapolated down to X-ray and UV/optical wavelengths."," Although we are working with a small sample of LAT bursts, and therefore suffer the consequences of small number statistics, our preliminary results indicate that LAT bursts exhibit bright late time X-ray and UV/optical afterglows because they are brighter at all wavelengths than the `average' burst, assuming the higher than average fluence can be extrapolated down to X-ray and UV/optical wavelengths."586" This seems to be the most likely explanation, given the known correlation between prompt emission and afterglow emission brightness (Gehrels et al."," This seems to be the most likely explanation, given the known correlation between prompt emission and afterglow emission brightness (Gehrels et al."587 2008)., 2008).588" We cannot say definitively, however, that this is the reason for the bright afterglows at late times, due to the presence of flares, which indicate possible late time central engine activity that could cause a rebrightening."," We cannot say definitively, however, that this is the reason for the bright afterglows at late times, due to the presence of flares, which indicate possible late time central engine activity that could cause a rebrightening."589" Without coverage of the early afterglow, it is impossible to say how the afterglow arrived at the state in which we observe it ~70 ks after the trigger."," Without coverage of the early afterglow, it is impossible to say how the afterglow arrived at the state in which we observe it $\sim$ 70 ks after the trigger."590" If we simply extrapolate the optical light curve of GRB 090926A backward, we find that they could have peaked as high as — 10 mag within the first hundred seconds after the trigger."," If we simply extrapolate the optical light curve of GRB 090926A backward, we find that they could have peaked as high as = 10 mag within the first hundred seconds after the trigger."591" Extrapolating the LAT spectrum of GRB 090926A to the band yields a peak magnitude of — 4, or if we assume a cooling break at GeV energies, the spectral index changes to DB5’—0.76, yielding a magnitude of = 15, consistent with our extrapolation backwards and the idea that LAT bursts are uniquely bright at all wavelengths."," Extrapolating the LAT spectrum of GRB 090926A to the band yields a peak magnitude of = 4, or if we assume a cooling break at GeV energies, the spectral index changes to $\beta \approx -0.76$, yielding a magnitude of = 15, consistent with our extrapolation backwards and the idea that LAT bursts are uniquely bright at all wavelengths."592" However, if the early afterglow was fainter than z 15 mag, then some sort of sustained energy injection would be required to keep the flux elevated at a level where we could then observe the bright afterglow at 70 ks after the trigger."," However, if the early afterglow was fainter than $\approx$ 15 mag, then some sort of sustained energy injection would be required to keep the flux elevated at a level where we could then observe the bright afterglow at 70 ks after the trigger."593 Such an energy injection would test our current theoretical understanding of GRB optical afterglows., Such an energy injection would test our current theoretical understanding of GRB optical afterglows.594 Our ability to determine the true nature of LAT-detected burst is contingent on our ability to follow-up LAT-detected GRBs at earlier times than has been achieved with the current sample., Our ability to determine the true nature of LAT-detected burst is contingent on our ability to follow-up LAT-detected GRBs at earlier times than has been achieved with the current sample.595"wwas detected with signal-Lo-noise ratio S/N>80 in each of its individual exposures: hence. the photometry is quite precise.BGo,","was detected with signal-to-noise ratio $S/N\geq80$ in each of its individual exposures; hence, the photometry is quite precise.,"596..ο. and wwere discovered with S/N in inclividual exposures (wpically 7.5. 2.7. and 2.4. respectively.," and were discovered with $S/N$ in individual exposures typically 7.5, 2.7, and 2.4, respectively."597" Discovery of the last object required the use of ""digital tracking. in which exposures are shifted at rates corresponding (to all valid NBO orbits before summing and searching for fIux peaks that exceed the detection threshold."," Discovery of the last object required the use of “digital tracking,” in which exposures are shifted at rates corresponding to all valid KBO orbits before summing and searching for flux peaks that exceed the detection threshold."598 Photometry for each object was extracted bv fitting a model of a moving point source to relevant expostres., Photometry for each object was extracted by fitting a model of a moving point source to relevant exposures.599 We measured the point spread function (PSF) lor the ACS Wide Field Camera (ΝΕΟ) in exposures of a globular cluster field., We measured the point spread function (PSF) for the ACS Wide Field Camera (WFC) in exposures of a globular cluster field.600 In the moving point source Chis PSF was smeared before fitting to the relevant pixels to account Dor (he (slight) (railing expected on each exposure.," In the moving point source model-fitting, this PSF was smeared before fitting to the relevant pixels to account for the (slight) trailing expected on each exposure."601 We fit (he entire stack of images simultaneously. with the free parameters being (he 6 relevant degrees of Ireedom in the NBO orbit plus an unknown [ας for each exposure.," We fit the entire stack of images simultaneously, with the free parameters being the 6 relevant degrees of freedom in the KBO orbit plus an unknown flux for each exposure."602 The best-fit photometry and orbit are tus solved simultaneously., The best-fit photometry and orbit are thus solved simultaneously.603 The S/N per exposure for iis so hieh that we use a slightly different approach. allowing the position to be a [ree parameter on each exposure rather (han forcing positions to obev a common orbit.," The $S/N$ per exposure for is so high that we use a slightly different approach, allowing the position to be a free parameter on each exposure rather than forcing positions to obey a common orbit."604 Without (his approach. we find (hat small (milliaresecond) errors in (he astrometric solutions for the Ες cause excess variance in the flux determinations.," Without this approach, we find that small (milliarcsecond) errors in the astrometric solutions for the WFC cause excess variance in the flux determinations."605 For the fainter three KABOs. the flux errors due to these ~5 milliaresecond astrometric errors are a few hundredthis of a magnitude. well below the noise levels.," For the fainter three KBOs, the flux errors due to these $\sim$ 5 milliarcsecond astrometric errors are a few hundredths of a magnitude, well below the noise levels."606 The slow brightening due to the decreasing illumination phase of ihe KBOs is too small to be detected in our data., The slow brightening due to the decreasing illumination phase of the KBOs is too small to be detected in our data.607 The fitting process produces uncertainties for each flux measurement., The fitting process produces uncertainties for each flux measurement.608 We find that the best-fit sinusoidal light curves give 4? per degree of [reedom (DOF) near unity for the three [aint bodies (see below). suggesting that our error estimates are reliable.," We find that the best-fit sinusoidal light curves give $\chi^2$ per degree of freedom (DOF) near unity for the three faint bodies (see below), suggesting that our error estimates are reliable."609 The 4? for the best sinusoidal fit for iis too high. partly because the light curve is clearly not sinusoidal (see below). but also because various svstemalic effects (e.g.. pointing jitter) mav allect (he PSF fitting at the 0.01 mag level.," The $\chi^2$ for the best sinusoidal fit for is too high, partly because the light curve is clearly not sinusoidal (see below), but also because various systematic effects (e.g., pointing jitter) may affect the PSF fitting at the 0.01 mag level."610 The formal errors on the magnitudes may also be underestimated. as is common for very hieh S/N photometry.," The formal errors on the magnitudes may also be underestimated, as is common for very high $S/N$ photometry."611 The midpoint Gime of each exposure is corrected for light-travel (ime trom the target., The midpoint time of each exposure is corrected for light-travel time from the target.612 The Gime-series photometry Lor (hese four objects is presented in Tables 1... 1., The time-series photometry for these four objects is presented in Tables \ref{bgtable} – \ref{fvtable}. .613 , 614"The magnetic field is mainly dissipated by the Ohmic loss when το,<| for main charged particles.",The magnetic field is mainly dissipated by the Ohmic loss when $|\tau_\nu \omega_\nu| < 1$ for main charged particles.615" La such a situation we have the approximate expression lrom equations (5)). (G)). C0)) and (3)) that where c, the electrical conductivity. q, aud im, are the electrical charge ancl the mass for a charged. particle v. respectively."," In such a situation we have the approximate expression from equations \ref{eq:diffusion_velocity}) ), \ref{eq:A}) ), \ref{eq:A1}) )and \ref{eq:A2}) ) that where $\sigma_{\rm c}$ the electrical conductivity, $q_\nu$ and $m_\nu$ are the electrical charge and the mass for a charged particle $\nu$ , respectively."616 Thus the drift velocity is independent of B., Thus the drift velocity is independent of $B$.617 On the other hand. when τω>1. the ambipolar diffusion is a main process of dissipation.," On the other hand, when $|\tau_\nu \omega_\nu| > 1$, the ambipolar diffusion is a main process of dissipation."618 In such a case we obtain where (he sullix i expresses (he ion particles., In such a case we obtain where the suffix i expresses the ion particles.619 Note that the drift velocity is proportional to D? in (his case., Note that the drift velocity is proportional to $B^2$ in this case.620 Tan&Blackman(2004) have evaluated the drift velocity of magnetic field in the quasi-hvdrostatie core found in Abel.Bryan&Norman(2000.2002).. when its density is a certain value.," \citet{TB04} have evaluated the drift velocity of magnetic field in the quasi-hydrostatic core found in \citet{abel00,abel02}, when its density is a certain value."621 However. (his svstem is highly non-equilibrium. and (he ionization degree eels smaller and smaller as (he collapse proceeds.," However, this system is highly non-equilibrium, and the ionization degree gets smaller and smaller as the collapse proceeds."622 Thus. its is not trivial at all whether the frozen condition is satisfied or not at much higher density.," Thus, its is not trivial at all whether the frozen condition is satisfied or not at much higher density."623 Therefore. we have to perform detailed non-ecuilibrium. caleulations of chemical reaction network. in order (o obtain the correct ionization degree in the collapsing gas.," Therefore, we have to perform detailed non-equilibrium calculations of chemical reaction network, in order to obtain the correct ionization degree in the collapsing gas."624 In order to investigate the evolution of ionized Iraction during the collapse in detail. we solve non-equilibrium chemical reaction network of primordial gas that involves not only 11 element. but also D. He. and Li.," In order to investigate the evolution of ionized fraction during the collapse in detail, we solve non-equilibrium chemical reaction network of primordial gas that involves not only H element, but also D, He, and Li."625 Furthermore.we introduce following24 species: e. I Η. HL. Ho. I5. EII. D. D.D. HD. UD. ID. Ile. He .Ie . Hell. Li. Li. Li Li .Li. Lill.and Lill .," Furthermore,we introduce following24 species: ${\rm e^-}$ ${\rm H^+}$ , H, ${\rm H^-}$, ${\rm H_2}$, ${\rm H_2^+}$ , ${\rm H_3^+}$, D, ${\rm D^+}$, $\rm D^-$, HD, ${\rm HD^+}$, ${\rm H_2D^+}$, He, ${\rm He^+}$ , ${\rm He^{++}}$ , ${\rm HeH^+}$, Li, ${\rm Li^+}$, ${\rm Li^{++}}$, ${\rm Li^{3+}}$, ${\rm Li^-}$, LiH,and ${\rm LiH^+}$ ."626 We emplov the latestreaction rate coellicients appearing in the following papers. Galli&Palla (1993).. Omukai (2000).. Stancil.Lepp&Dalearno(1998 )..Flower(2002) and Lepp.Stancil&Dalgarno (2002).," We employ the latestreaction rate coefficients appearing in the following papers, \citet{galli98}, \citet{omukai00}, , \citet{stancil98}, ,\citet{flower02} and \citet{lepp02}."627.As [orthe radiative recombination. we use (he rate coefficients basedon Spitzer(LOTS).," .As forthe radiative recombination, we use the rate coefficients basedon \citet{spitzer78}.."628The spatial distribution of the candidates appear to be in a filamentary structure. with a 4o confidence. however to confirm this and to plot the filament in 3D-space. we would need spectroscopic redshifts.,"The spatial distribution of the candidates appear to be in a filamentary structure, with a $\sigma$ confidence, however to confirm this and to plot the filament in 3D-space, we would need spectroscopic redshifts."629 We have studied the entire candidate sample m all bands available from X-rays to infrared in the GOODS-S data-set., We have studied the entire candidate sample in all bands available from X-rays to infrared in the GOODS-S data-set.630 From the SED fitting we conclude that the LEGOs on average have low metallicity (Z/Z..= 0.005). have stellar masses in the range of 15«10 M. and low dust extinction (Ay-0.9).," From the SED fitting we conclude that the LEGOs on average have low metallicity $Z/Z_{\odot} = 0.005$ ), have stellar masses in the range of $1 - 5 \times 10^9$ $_{\odot}$ and low dust extinction $_V \sim 0.3$ )."631 The candidates have ages in the range of 100 — 900 Myrs., The candidates have ages in the range of 100 – 900 Myrs.632 We also find one galaxy. LEGO_GGOODS-S#116. which is best fit by a dusty starburst galaxy at 2=3.15 with Lya-emission escaping from an area slightly offset from the central core.," We also find one galaxy, 16, which is best fit by a dusty starburst galaxy at $z = 3.15$ with $\alpha$ -emission escaping from an area slightly offset from the central core."633 The comparison to à sample of ('-band drop-out galaxies in the GOODS-S field show that the colours of LEGOs are consistent with the selection criteria for ÜC-band. drop-outs except they are too faint to be detected from their continuum flux., The comparison to a sample of $U$ -band drop-out galaxies in the GOODS-S field show that the colours of LEGOs are consistent with the selection criteria for $U$ -band drop-outs except they are too faint to be detected from their continuum flux.634 They also have colours similar to those of LBGs at redshift z3., They also have colours similar to those of LBGs at redshift $z \approx 3$.635 In agreement with previous results (e.g. Gawiser et al., In agreement with previous results (e.g. Gawiser et al.636 2006). we conclude that Lyo-emitters at redshift .~3.1 are dust- and AGN-free. star-forming galaxies with small to medium masses.," 2006), we conclude that $\alpha$ -emitters at redshift $z \sim 3.1$ are dust- and AGN-free, star-forming galaxies with small to medium masses."637The McDonald spectrum for the abundance analysis was obtained on the night of 2008 August 10 (JD 2454688.7) at the 2.7 meter Harlan J. Smith reflector with the Tull cross-dispersed éechelle spectrograph (Tull et al.,The McDonald spectrum for the abundance analysis was obtained on the night of 2008 August 10 (JD 2454688.7) at the 2.7 meter Harlan J. Smith reflector with the Tull cross-dispersed écchelle spectrograph (Tull et al.638 1995) at à spectral resolution of 60. 000.," 1995) at a spectral resolution of $\lambda/d\lambda \simeq 60,000$ ."639 The spectrum covers the wavelength ranges 3800 tto 105500 wwith no gaps in the wavelength ranges 3800 tto 4885 aand 5020 tto 5685 but coverage is incomplete but substantial beyond 5700À:: the effective short and long wavelength limits are set by the useful S/N ratio., The spectrum covers the wavelength ranges 3800 to 500 with no gaps in the wavelength ranges 3800 to 4885 and 5020 to 5685 but coverage is incomplete but substantial beyond 5700; the effective short and long wavelength limits are set by the useful S/N ratio.640 A ThAr hollow cathode lamp provided the wavelength calibration., A ThAr hollow cathode lamp provided the wavelength calibration.641 Flat-field and bias exposures completed the calibration files., Flat-field and bias exposures completed the calibration files.642 The signal-to-noise ratio ranges between 40 and 90 per pixel. not only changing with the blaze function within echelle orders but also star brightness between echelle orders.," The signal-to-noise ratio ranges between 40 and 90 per pixel, not only changing with the blaze function within echelle orders but also star brightness between echelle orders."643 The MeD observations were reduced using the echelle reduction package (Mills Webb 1994)., The McD observations were reduced using the echelle reduction package (Mills Webb 1994).644 The spectra were extracted using implementation of the optimal extraction algorithm developed by Horne (1986)., The spectra were extracted using implementation of the optimal extraction algorithm developed by Horne (1986).645 propagates error. information based on. photon statistics and readout noise throughout the extraction process., propagates error information based on photon statistics and readout noise throughout the extraction process.646 The bias level in the overscan area was modeled with a polynomial and subtracted., The bias level in the overscan area was modeled with a polynomial and subtracted.647 The scattered light was modeled and removed from the spectrum., The scattered light was modeled and removed from the spectrum.648 In order to correct for pixel-to-pixel sensitivity variations. “flattield” exposures from a halogen lamp were used.," In order to correct for pixel-to-pixel sensitivity variations, `flatfield' exposures from a halogen lamp were used."649 Individual orders were cosmic-ray cleaned. and continuum normalized with bespoke echelle reduction software in (Saahin. 2008).," Individual orders were cosmic-ray cleaned, and continuum normalized with bespoke echelle reduction software in (Şaahin 2008)."650 Reduced spectra were transferred to the spectrum analysis program (Howarth et al., Reduced spectra were transferred to the spectrum analysis program (Howarth et al.651 1998) for further analysis (e.g. for equivalent width measurement)., 1998) for further analysis (e.g. for equivalent width measurement).652 In equivalent width measurements. local continua on both side of the lines were fitted with a first-degree polynomial then equivalent widths were measured with respect to these local continua using a fitted Gaussian profile.," In equivalent width measurements, local continua on both side of the lines were fitted with a first-degree polynomial then equivalent widths were measured with respect to these local continua using a fitted Gaussian profile."653 For strong lines. a direct integration was preferred to the Gaussian approximation.," For strong lines, a direct integration was preferred to the Gaussian approximation."654 The errors for each equivalent width measurement were determined on the basis of scatter of linear continuum fit and signal-to-noise ratio of each measured line in the spectra., The errors for each equivalent width measurement were determined on the basis of scatter of linear continuum fit and signal-to-noise ratio of each measured line in the spectra.655 Errors on the measured equivalent widths are calculated using the prescriptions given by Howarth Phillips €1986)., Errors on the measured equivalent widths are calculated using the prescriptions given by Howarth Phillips (1986).656 0.2 em The SAO spectrum was obtained on the night of 2009 June 10 (JD 2452993.4) by VK and NST with the NES echelle spectrograph mounted at the Nasmyth focus of the 6-m telescope of the Special Astrophysical Observatory (Panchuk et al., 0.2 cm The SAO spectrum was obtained on the night of 2009 June 10 (JD 2452993.4) by VK and NST with the NES echelle spectrograph mounted at the Nasmyth focus of the 6-m telescope of the Special Astrophysical Observatory (Panchuk et al.657 2007) with a 2048 X 2048 CCD with an image slicer (Panchuk et al., 2007) with a 2048 X 2048 CCD with an image slicer (Panchuk et al.658 2007) and a spectral resolution of αλz60.000.," 2007) and a spectral resolution of $\lambda/d\lambda \ge 60,000$."659 A modified ECHELLE context (Yushkin Klochkova 2005) of package was used to extract one-dimensional vectors from the two-dimensional echelle spectra., A modified ECHELLE context (Yushkin Klochkova 2005) of package was used to extract one-dimensional vectors from the two-dimensional echelle spectra.660 Wavelength calibration was performed using a hollow-cathode Th-Ar lamp., Wavelength calibration was performed using a hollow-cathode Th-Ar lamp.661 The wavelength coverage for the SAO spectrum was +460 — 5920., The wavelength coverage for the SAO spectrum was 4460 – 5920.662 Measurement of equivalent widths was carried out as for the McDonald spectrum., Measurement of equivalent widths was carried out as for the McDonald spectrum.663 (0.2 em The agreement between the McDonald and SAO spectra is satisfactory. i.e.. the two spectra are similar as to line width. depth and equivalent width for weak to strong lines.," 0.2 cm The agreement between the McDonald and SAO spectra is satisfactory, i.e., the two spectra are similar as to line width, depth and equivalent width for weak to strong lines."664 This is shown by the section of the reduced spectra illustrated in Figure |.., This is shown by the section of the reduced spectra illustrated in Figure \ref{f_spectrum_part}.665 A sample of apparently unblended lines was selected from across the common wavelength interval and their equivalent widths (EWs) measured in both the McDonald and the SAO spectra., A sample of apparently unblended lines was selected from across the common wavelength interval and their equivalent widths (EWs) measured in both the McDonald and the SAO spectra.666 The comparison of EWS shown in Figure 2. shows good agreement between the two sets of measurements., The comparison of EWs shown in Figure \ref{f_ew_comparison} shows good agreement between the two sets of measurements.667 Across the common wavelength interval. we compare EWs from MeDonald and SAO spectra. especially for lines at the limit of detection and for elements represented by just one or two lines.," Across the common wavelength interval, we compare EWs from McDonald and SAO spectra, especially for lines at the limit of detection and for elements represented by just one or two lines."668 The SAO spectrum was used to provide lines that fell in the inter-order gaps of the McDonald spectrum., The SAO spectrum was used to provide lines that fell in the inter-order gaps of the McDonald spectrum.669 Although the data are sparse. the star is probably not a large amplitude velocity variable.," Although the data are sparse, the star is probably not a large amplitude velocity variable."670 Klochkova (1995) reported a heliocentric radial velocity οἳ 32.5+0.4kms 1., Klochkova (1995) reported a heliocentric radial velocity of $-$ $\pm$ 0.4 km $^{-1}$.671" Hrivnak. Kwok. Volk (1988) report 30+2 kms + from an unspecitied number of measurements but add that there is ""an indication of variability’."," Hrivnak, Kwok, Volk (1988) report $-$ $\pm$ 2 km $^{-1}$ from an unspecified number of measurements but add that there is `an indication of variability'."672 The McD spectrum gives 3041kms ! from the metal lines., The McD spectrum gives $-$ $\pm$ 1 km $^{-1}$ from the metal lines.673 The 2009 SAO spectrum gives 31.8+1.7 kms +., The 2009 SAO spectrum gives $-$ $\pm$ 1.7 km $^{-1}$.674 Although Lewis. Eder. Terzian (1985) and Eder. Lewis. Terzian (1983) cite the heliocentric velocity as — 17.4 km + from their observed OH radio lines. Hrivnak (2009. private communication) indicates that this value resulted from an incorrect conversion of LSR to heliocentric velocity and a velocity of about 30 km 5 is obtained from the OH velocities.," Although Lewis, Eder, Terzian (1985) and Eder, Lewis, Terzian (1988) cite the heliocentric velocity as $-$ 17.4 km $^{-1}$ from their observed OH radio lines, Hrivnak (2009, private communication) indicates that this value resulted from an incorrect conversion of LSR to heliocentric velocity and a velocity of about $-30$ km $^{-1}$ is obtained from the OH velocities."675 (0.2 em The Na D lines show four components., 0.2 cm The Na D lines show four components.676 Figure 3. shows the Na D from the MeDonald and SAO spectra.," Figure \ref{NaD_lines}677 shows the Na D from the McDonald and SAO spectra."678 Heliocentric velocities of the four principal components in the McDonald spectrum are isted in Table |., Heliocentric velocities of the four principal components in the McDonald spectrum are listed in Table 1.679 The SAO spectrum gives similar velocities., The SAO spectrum gives similar velocities.680 Stellar photospherie Na lines atabout 30km + must be largely masked by these multiple circumstellar components., Stellar photospheric Na lines at about $-$ 30 km $^{-1}$ must be largely masked by these multiple circumstellar components.681 Component |. if not an interstellar component. represents an outflow at a velocity of about I2 kms. +.," Component 1, if not an interstellar component, represents an outflow at a velocity of about 12 km $^{-1}$."682 Components 3 and 4 are falling toward the star at velocities of about [2 and 25 kms. +. respectively.," Components 3 and 4 are falling toward the star at velocities of about 12 and 25 km $^{-1}$, respectively."683 (0.2 em Weak emission in the blue and red wings of the Ha profile flanking a deep narrow absorption core was reported by Klochkova (1995 - see also Tamura. Takeuti Zalewski 1993).," 0.2 cm Weak emission in the blue and red wings of the $\alpha$ profile flanking a deep narrow absorption core was reported by Klochkova (1995 - see also Tamura, Takeuti Zalewski 1993)."684 On the McDonald spectrum. Ha occurs at the very edge of an order but à very shallow a deep core flanked by red emission is seen.," On the McDonald spectrum, $\alpha$ occurs at the very edge of an order but a very shallow a deep core flanked by red emission is seen."685 H.? and higher lines in the Balmer series and Paschen lines are purely in absorption., $\beta$ and higher lines in the Balmer series and Paschen lines are purely in absorption.686 The 2009 SAO spectrum did not include Ha., The 2009 SAO spectrum did not include $\alpha$ .687 (0.2 em The stellar absorption lines are broad., 0.2 cm The stellar absorption lines are broad.688 If one accepts. classical notions of microturbulence and maeroturbulenee. this width suggests substantial macroturbulence in the atmosphere.," If one accepts classical notions of microturbulence and macroturbulence, this width suggests substantial macroturbulence in the atmosphere."689 The microturbulence is about5 km + (see below)., The microturbulence is about5 km $^{-1}$ (see below).690 The instrumental, The instrumental691detailed PIC simulations. however. the filament evolution. the isotropization of the upstream particles. and the production of shock jump conditions remain poorly understood.,"detailed PIC simulations, however, the filament evolution, the isotropization of the upstream particles, and the production of shock jump conditions remain poorly understood."692 These processes. of course. are of Immense Importance for the microphysics of GRBs and collisionless shocks in general.," These processes, of course, are of immense importance for the microphysics of GRBs and collisionless shocks in general."693 The two main processes that dictate the evolution the ordered outcome of the transverse Weibel instability are the magnetic self-confinement of each filament and the interactions. between. neighboring. filaments., The two main processes that dictate the evolution the ordered outcome of the transverse Weibel instability are the magnetic self-confinement of each filament and the interactions between neighboring filaments.694 Here. we use simple analytical arguments to explore the first one. namely the stability of a single filament.," Here, we use simple analytical arguments to explore the first one, namely the stability of a single filament."695 In 2.. we present an MHD model of such current filament.," In \ref{sec:equilibrium}, we present an MHD model of such current filament."696 In .3.. we show that the filament is unstable to a kink-like mode. and estimate the growth rate of the mode.," In \ref{sec:stability}, we show that the filament is unstable to a kink-like mode, and estimate the growth rate of the mode."697 This instability destroys the quasitwo-dimensional geometry of the filament and produces a significant magnetic field in the direction parallel to shock propagation., This instability destroys the quasi--two-dimensional geometry of the filament and produces a significant magnetic field in the direction parallel to shock propagation.698 In 4.. we study the motion of collisionless test particles in the background of a current filament undergoing the instability.," In \ref{sec:transport}, we study the motion of collisionless test particles in the background of a current filament undergoing the instability."699 We show that charged particles confined within the filament diffuse in energy space., We show that charged particles confined within the filament diffuse in energy space.700 We suggest that this diffusion tsotropisizes and thermalizes the upstream particles., We suggest that this diffusion isotropisizes and thermalizes the upstream particles.701 In 6.. we discuss implications of these processes for the magnetic field decay and the physics of GRB afterglows.," In \ref{sec:discussion}, we discuss implications of these processes for the magnetic field decay and the physics of GRB afterglows."702 Consider a magnetostatic equilibrium composed of an infinite cylindrical current filament extending in the cz direction., Consider a magnetostatic equilibrium composed of an infinite cylindrical current filament extending in the $z$ direction.703 Let r=0 be the center of the filament. where we use cylindrical coordinates (70.2).," Let $r=0$ be the center of the filament, where we use cylindrical coordinates $(r,\theta,z)$."704 Let the current flow in the +: direction along the central axis., Let the current flow in the $+z$ direction along the central axis.705 Further. assume that at some radius r=R. the current drops to zero.," Further, assume that at some radius $r=R$, the current drops to zero."706 We refer to R as the radius of the filament., We refer to $R$ as the radius of the filament.707 At r>R. there are currents flowing in other filaments in both directions (+2). so that the total current flowing through any given plane z2const vanishes.," At $r>R$, there are currents flowing in other filaments in both directions $\pm z$ ), so that the total current flowing through any given plane $z={\rm708const}$ vanishes."709 Since we are interested in the stability of a single filament. we ignore the effect on the filament of currents external to the filament.," Since we are interested in the stability of a single filament, we ignore the effect on the filament of currents external to the filament."710 Particles giving rise to the current must move coherently. e.g.. If there is a pinching azimuthal magnetic field By> 0. positive charges move in the positive z-direction.. and negative charges move in the negative direction.," Particles giving rise to the current must move coherently, e.g., if there is a pinching azimuthal magnetic field $B_\theta>0$ , positive charges move in the positive $z$ -direction, and negative charges move in the negative $z$ -direction."711 This is possible if the particles do not execute a full gyration in the magnetic field (Alfvén1939):: fractional gyrations give rise to directed current as explained in Spitzer(1965)., This is possible if the particles do not execute a full gyration in the magnetic field \citep{Alfven:39}; fractional gyrations give rise to directed current as explained in \citet{Spitzer:65}.712. This requirement implies that a filament’s radial structure depends on its radius., This requirement implies that a filament's radial structure depends on its radius.713 A detailed discussion of this structure can be found in Davidson(1974) and Honda(200001: we here summarize the relevant aspects., A detailed discussion of this structure can be found in \citet{Davidson:74} and \citet{Honda:00}; we here summarize the relevant aspects.714" Consider first the regime in which the radius of the filament R is smaller than about the plasma skin depth. R© 6. The skin depth equals 6=o>!""wp. Where wy=(Axenim)L2 is the plasma frequency. 5 is the kinetic Lorentz factor the particles. wis the mass of the particles. and 71 is their density (in e shocks. z is the electron. mass: in ep shocks. it is the proton mass)."," Consider first the regime in which the radius of the filament $R$ is smaller than about the plasma skin depth, $R \lesssim \delta$ The skin depth equals $\delta\equiv c\gamma^{1/2}/\omega_{\rm715p}$, where $\omega_{\rm p}\equiv (4\pi e^2 n/m)^{1/2}$ is the plasma frequency, $\gamma$ is the kinetic Lorentz factor the particles, $m$ is the mass of the particles, and $n$ is their density (in $e^\pm$ shocks, $m$ is the electron mass; in $e^-p$ shocks, it is the proton mass)."716" This directly implies that the current flowing through the filament /~zRne where is the average axial velocity of particles in the fic,filament. Jicdoes not exceed the Alfvénn critical current Z4—Hunce."," This directly implies that the current flowing through the filament $I717\sim \pi R^2 ne\beta_\parallel c$, where $\beta_\parallel c$ is the average axial velocity of particles in the filament, does not exceed the Alfvénn critical current $I_{\rm A} =718\gamma\beta_\parallel m c^3/e$."719 Since the magnetic field is related to the current via By~7/Re. the condition /=/4 implies that the Larmor radius of the particle n25mc/eBg must exceed the radius of the filament m2(3/3)R>ΝΔ.," Since the magnetic field is related to the current via $B_\theta \sim I/Rc$, the condition $I\lesssim I_{\rm A}$ implies that the Larmor radius of the particle $r_{\rm L}=\gamma \beta m c^2 / e B_\theta$ must exceed the radius of the filament $ r_{\rm L} \gtrsim720(\beta/\beta_\parallel) R > R$."721 As a result. when Rz6. all particles magnetically confined. to the filament move in directed fashion.," As a result, when $R \lesssim \delta$, all particles magnetically confined to the filament move in directed fashion."722" Here 7, depends weakly on r and thus the current profile across the filament is approximately homogeneous.", Here $\beta_\parallel$ depends weakly on $r$ and thus the current profile across the filament is approximately homogeneous.723 However when R>6. if the current were uniform within r=R and the average axial velocity of the current carrying charges were relativistic (3)~D). the Larmor radius would be smaller than the radius of the filament. and directed motion of charges within the filament would be compromised.," However when $R > \delta$, if the current were uniform within $r\lesssim R$ and the average axial velocity of the current carrying charges were relativistic $\beta_\parallel\sim 1$ ), the Larmor radius would be smaller than the radius of the filament, and directed motion of charges within the filament would be compromised."724 Directed current flow can still be maintained when R>6 if the current is confined within a thin annular cylindrical region of width AR not exceeding the Larmor radius., Directed current flow can still be maintained when $R>\delta$ if the current is confined within a thin annular cylindrical region of width $\Delta R$ not exceeding the Larmor radius.725 The current then gives rise to a thin magnetic “wall” against which a particle confined to the filament can be reflected., The current then gives rise to a thin magnetic “wall” against which a particle confined to the filament can be reflected.726" Then the Larmor radius within the current carrying layer can be written r=V/AR. and the condition rj,zAR implies that the current carrying layer is thinner than the skin depth. ARὁ, as seen in PIC simulations."," Then the Larmor radius within the current carrying layer can be written $r_{\rm L}=\delta^2/\Delta R$, and the condition $r_{\rm L}\geq\Delta R$ implies that the current carrying layer is thinner than the skin depth, $\Delta R\leq \delta$, as seen in PIC simulations."727 In this regime. the saturated phase of the transverse Weibel instability consists current-carrying domains separated by thin magnetic walls.," In this regime, the saturated phase of the transverse Weibel instability consists current-carrying domains separated by thin magnetic walls."728 Analytic considerations (Milosavljevic.aka|&Spitkovsky2005) απά numerical— simulations indicate that once the transverse Weibel instability saturates in collisionless shocks. the filament size 1s comparable to the skin depth.," Analytic considerations \citep{Milosavljevic:05} and numerical simulations indicate that once the transverse Weibel instability saturates in collisionless shocks, the filament size is comparable to the skin depth."729 Motivated by these results we focus here on the case in which R—ó., Motivated by these results we focus here on the case in which $R \sim \delta$.730 We assume that MHD equations apply and that the pressure tensor of the particles is isotropic., We assume that MHD equations apply and that the pressure tensor of the particles is isotropic.731 Both of these assumptions are oversimplifications., Both of these assumptions are oversimplifications.732 The pressure tensor near the axis of the filament is approximately isotropic when rp~ à.," The pressure tensor near the axis of the filament is approximately isotropic when $r_{\rm L}\sim R \sim733\delta$ ."734 The purpose of the MHD model is to elucidate the physical mechanisms and motivate specific collisionless PIC simulations of the saturated state of the Weibel instability., The purpose of the MHD model is to elucidate the physical mechanisms and motivate specific collisionless PIC simulations of the saturated state of the Weibel instability.735 The simulations are the best way to test the theory over a range of parameter values., The simulations are the best way to test the theory over a range of parameter values.736 The toroidal magnetic field B=BG) is related to the axial current density J2(6702 via J2l———«rixa," The toroidal magnetic field ${\bf737B}=B(r){\bf \hat \theta}$ is related to the axial current density ${\bf J}=J(r) {\bf \hat z}$ via J = (rB)."738 Within the MHD approximation. the fluid pressure 01) satisfies the equation of pressure equilibrium VP=J\B," Within the MHD approximation, the fluid pressure $P(r)$ satisfies the equation of pressure equilibrium P = = B^2)."739 To construct an magnetostatic equilibrium filament. one can choose the radial dependence of the magnetic field. and then evaluate the current density and the pressure using equations (2)) and (2)).," To construct an magnetostatic equilibrium filament, one can choose the radial dependence of the magnetic field, and then evaluate the current density and the pressure using equations \ref{eq:current_density}) ) and \ref{eq:grad_pressure}) )."740 It immediately follows from equation (2)) that the fluid pressure inside the filament is larger than outside. the magnetic pressure accounting for the difference.," It immediately follows from equation \ref{eq:grad_pressure}) ) that the fluid pressure inside the filament is larger than outside, the magnetic pressure accounting for the difference."741 This pressure imbalance is the origin of the unstable behavior that we explore below., This pressure imbalance is the origin of the unstable behavior that we explore below.742 The filaments are unstable to the well-known sausage. kink. and related MHD modes (see..," The filaments are unstable to the well-known sausage, kink, and related MHD modes (see.,"743 e.g.. Hasegawa and references therein) which result in the distortion of the filament boundary.," e.g., \citealt{Hasegawa:75} and references therein) which result in the distortion of the filament boundary."744We here focus on a particular mode. the helical kink instability. but expect similar stability criteria. growth rates. and particle transport in other related modes.,"We here focus on a particular mode, the helical kink instability, but expect similar stability criteria, growth rates, and particle transport in other related modes."745 Consider linear magnetostatic perturbations around the equilibrium described in 2.., Consider linear magnetostatic perturbations around the equilibrium described in \ref{sec:equilibrium}. .746 According to the energy principle. elegantly proven in Kulsrud (2005).. a perturbation," According to the energy principle, elegantly proven in \citet{Kulsrud:05}, , a perturbation"747"For redshifts in the range 0.71.0. we find. acceptable DUSTY models with 7=55SO (equivalent to the fitted Ny=69.17cm for a Galactic dust:gas ratio) for R=100 απ 7i,=1500 Ix: they have ri,20.55.L0 and Qus&1.0.","For redshifts in the range 0.7–1.0, we find acceptable DUSTY models with $\tau=55-80$ (equivalent to the fitted $N_{\rm{H}}=6-9 \times 10^{22}$ for a Galactic dust:gas ratio) for $R=100$ and $T_{\rm{in}}=1500$ ; they have $r_{\rm{in}}=0.55-1.0$ and $\alpha_{\rm{ox}} \simeq 1.0$."748 ]t is instructive to compare our results with the calculations of Granato. Danese Franceschini (1997). who compared dust radiative transfer models with the infrared properties of Sevlert galaxies.," It is instructive to compare our results with the calculations of Granato, Danese Franceschini (1997), who compared dust radiative transfer models with the infrared properties of Seyfert galaxies."749 They concluded that the moderately thick. extended tori of GCranato Danese (1994) (GD) (with 5—ely<SO mag and outer radii of 10's to 100's of parsec). provided a better fit to the data than the thick. very compact models of Pier Ixrolik (1992a) (PIN) (with lycSOO mae. all within a pc)).," They concluded that the moderately thick, extended tori of Granato Danese (1994) (GD) (with $5 \leq A_{\rm{V}} \leq 80$ mag and outer radii of 10's to 100's of parsec), provided a better fit to the data than the thick, very compact models of Pier Krolik (1992a) (PK) (with $A_{\rm{V}} \geq 800$ mag, all within a )."750 They also found thatthe observed. X-ray, They also found thatthe observed X-ray751Dwarf galaxies are the galaxies with lower luminosity. lower mass and smaller size. and they are much more abundant in number than normal galaxies on the Hubble sequence like our Milky Way (seeKunth&Ostlin2000.forarecentreviewoflocalgroupdwarf galaxies)...,"Dwarf galaxies are the galaxies with lower luminosity, lower mass and smaller size, and they are much more abundant in number than normal galaxies on the Hubble sequence like our Milky Way \citep[see][for a752 recent review of local group dwarf galaxies]{ko00}."753 These natures imply that they are basic building blocks of the normal galaxies in hierarchical structure formation scenarios of the cold dark matter universe (Coleetal.2000;Corbin&Vacea2002;Monica2003).," These natures imply that they are basic building blocks of the normal galaxies in hierarchical structure formation scenarios of the cold dark matter universe \citep{clbf00, cv02, mo03}."754. In addition. the dwarf galaxies have lower abundances of heavy elements than those of the normal galaxies: therefore. they are still in an early stage of chemical evolution and/or are analogies to primordial galaxies (e.g..Izotov&Thuan1999).," In addition, the dwarf galaxies have lower abundances of heavy elements than those of the normal galaxies; therefore, they are still in an early stage of chemical evolution and/or are analogies to primordial galaxies \citep[e.g.,][]{it99}."755. Dwarf galaxies. are classified. by their norphological appearances into some basic types., Dwarf galaxies are classified by their morphological appearances into some basic types.756 One of the morphological types is dwarf irregular galaxies (dlrs). which characteristics are irregular and amorphous appearances at optical wavelengths.," One of the morphological types is dwarf irregular galaxies (dIrrs), which characteristics are irregular and amorphous appearances at optical wavelengths."757 Interestingly. unlike other norphological types of dwarf galaxies. dlrrs still have plentiful interstellar medium (SM) and 1n general show ongoing star formation and HII region (Mateo1998:Kunth&Ostlin2000).," Interestingly, unlike other morphological types of dwarf galaxies, dIrrs still have plentiful interstellar medium (ISM) and in general show ongoing star formation and HII region \citep{ma98, ko00}."758.. The star formation histories (SFHs) of dIrrs are considered to be episodic. that is. quiescent phase follows a starburst phase (Greggioetal.1993:Tolstoy1998:Grebel1999).. and their star formation rates (SFRs) are more widely distributed ἐς107=1077[Mo/yr/kpc p than that of the normal galaxies (Hunter1997).," The star formation histories (SFHs) of dIrrs are considered to be episodic, that is, quiescent phase follows a starburst phase \citep{greggio93, tolstoy98, grebel99}, and their star formation rates (SFRs) are more widely distributed $\sim759 10^{-4}-10^{-2}\dsfr$ ) than that of the normal galaxies \citep{hunter97}."760. In this issue. we particularly focus on the dIrrs which are considered to be in their quiescent phase of star formation (i.e. SER~107 [Ma/yr/kpe?}).," In this issue, we particularly focus on the dIrrs which are considered to be in their quiescent phase of star formation (i.e. $\sfr \sim 10^{-4}\dsfr$ )."761 However. despite their closeness. the regulation of star formation even in the Local Group dlrrs ts not yet well understood.," However, despite their closeness, the regulation of star formation even in the Local Group dIrrs is not yet well understood."762" Actually. it is very surprising that they have been forming stars recently because their gravitational potential are shallow (Skillman&Bender1997.forareview) and comparable to the energies of some supernova explosions: therefore. they should have little ISM after some supernovae (several million years from the epoch of dirrs formed). and thus. the recent and ongoing star formation would not occur,"," Actually, it is very surprising that they have been forming stars recently because their gravitational potential are shallow \citep[][for a review]{sb97}763 and comparable to the energies of some supernova explosions; therefore, they should have little ISM after some supernovae (several million years from the epoch of dIrrs formed), and thus, the recent and ongoing star formation would not occur."764 Why do they still have plentiful gas?, Why do they still have plentiful gas?765 Why do they show ongoing star formation?, Why do they show ongoing star formation?766 Both the questions are. still controversial., Both the questions are still controversial.767 Even if these questions are solved. another puzzle remains; why do they form stars at such low rates in spite of having plentiful ISM as material of stars?," Even if these questions are solved, another puzzle remains; why do they form stars at such low rates in spite of having plentiful ISM as material of stars?"768 This paper attempts to solve this puzzle by investigating the notchemical equilibrium molecule formation history of dlrrs o=) some assumptions., This paper attempts to solve this puzzle by investigating the non-chemical equilibrium molecule formation history of dIrrs on some assumptions.769 For the star formation processes in low-metallicity environment. hydrogen molecules (Hs) can be important coolants at the gas temperature T€10K (Peebles 1968):: more abundant leads gas to cool faster. and therefore. would result in more active star formation.," For the star formation processes in low-metallicity environment, hydrogen molecules $\mathrm{H_2}$ ) can be important coolants at the gas temperature $T\la 10^4\dkel$ \citep{pd68}; more abundant leads gas to cool faster, and therefore, would result in more active star formation."770 Although forms primarily on the surfaces of dust grains in local molecular clouds. can also form in gas-phase through the reactions. which intermediaries are negatively charged hydrogen (H). or ionized hydrogen molecule (H5 ). These gas-phase reactions are active T~10K and the reactions (1) and (2) can more efficiently form than the (3) and (4) reactions (e.g..Abeletal.1997).," Although forms primarily on the surfaces of dust grains in local molecular clouds, can also form in gas-phase through the reactions, which intermediaries are negatively charged hydrogen $\mathrm{H^-}$ ), or ionized hydrogen molecule $\mathrm{H_2^+}$ ), These gas-phase reactions are active $T\sim 10^4\dkel$ and the reactions ) and ) can more efficiently form than the ) and ) reactions \citep[e.g.,][]{abel97}."771. On the other hand. ts dissociated by collisions with other chemical species and the absorption of ultraviolet (UV) photons called the Lyman-Werner band (LW) photons (11.26eV</iv€13.6 eV). which are mainly radiated by early type stars (Field.Somerville&Dressler 1966).," On the other hand, is dissociated by collisions with other chemical species and the absorption of ultraviolet (UV) photons called the Lyman-Werner band (LW) photons $11.26\ \mathrm{eV} \le h\nu \le 13.6\ \mathrm{eV}$ ), which are mainly radiated by early type stars \citep{fsd66}."772. Many astrophysicists have built the theoretical. models concentrated on the dissociation of by the LW photons because they have investigated. the second epoch of star formation after the first stars have been formed in metal-free molecular cloud (e.g..Silk1977;Nishi&Tashiro2000).," Many astrophysicists have built the theoretical models concentrated on the dissociation of by the LW photons because they have investigated the second epoch of star formation after the first stars have been formed in metal-free molecular cloud \citep[e.g.,][]{silk77,773 nt00}."774. These models result in the low abundance of in relatively large region (several hundred parsec around the source star of the LW photons). and thus. negative-feedback on further star formation in the host cloud of the star.," These models result in the low abundance of in relatively large region (several hundred parsec around the source star of the LW photons), and thus, negative-feedback on further star formation in the host cloud of the star."775 However. in dlrrs. the region affected by LW photons would be smaller than previously predicted because dlrrs are gas rich system: the mean free paths of LW photons are on the order of 1pe (« the typical," However, in dIrrs, the region affected by LW photons would be smaller than previously predicted because dIrrs are gas rich system; the mean free paths of LW photons are on the order of $1\ \mathrm{pc}$ $\ll$ the typical"776I. rather than a change in the absorption column density or in the properties of the soft component.,"$\Gamma$, rather than a change in the absorption column density or in the properties of the soft component."777 Indeed. the value of Nyy. as well as the temperature and normalization of the BB and the MKL components. are compatible with being constant (to within the errors) in the three spectra.," Indeed, the value of $N_{\rm H}$, as well as the temperature and normalization of the BB and the MKL components, are compatible with being constant (to within the errors) in the three spectra."778 That most of the changes in the X-ray flux of the source occurred in the 1.5- keV energy band. where the contribution of the CUTOFFPL component is much greater than that of the soft component (see Fig. 9)).," That most of the changes in the X-ray flux of the source occurred in the 1.5-10 keV energy band, where the contribution of the CUTOFFPL component is much greater than that of the soft component (see Fig. \ref{fig:xtesp1unfolded}) ),"779" added support fot the above conclusion,", added support fot the above conclusion.780 To further constrain the soft component more tightly. we also extracted the source spectrum by using the total available exposure time of the oobservation.," To further constrain the soft component more tightly, we also extracted the source spectrum by using the total available exposure time of the observation."781 A fit to this spectrum. with a simple BB or a PL model did not provide an acceptable result (L2 Alo.f24.9/101. 18/101. respectively for the BB and the PL model).," A fit to this spectrum with a simple BB or a PL model did not provide an acceptable result $\chi^2_{\rm red}$ /d.o.f=4.9/101, 1.8/101, respectively for the BB and the PL model)."782" A CUTOFFPL model only marginally improved the fit (\2,,/d.0.f=1.5/100) and some structures. remained present in the residuals (see Fig. 13)).", A CUTOFFPL model only marginally improved the fit $\chi^2_{\rm red}$ /d.o.f=1.5/100) and some structures remained present in the residuals (see Fig. \ref{fig:igrbsp}) ).783 To obtain an acceptable fit to the data. we thus used the same spectral models that we adopted for the rate-resolved analysis.," To obtain an acceptable fit to the data, we thus used the same spectral models that we adopted for the rate-resolved analysis."784 All these models provided an equivalently good fit to the data., All these models provided an equivalently good fit to the data.785 The results of these fits are given in Table 4+ and discussed in detail in the following section., The results of these fits are given in Table \ref{tab:igrbfittotal} and discussed in detail in the following section.786 We note that when fitting the total spectrum of wwe discarded data in the energy range 0.4-0.6 keV. as we noted that in this energy range the background was rather high.," We note that when fitting the total spectrum of we discarded data in the energy range 0.4-0.6 keV, as we noted that in this energy range the background was rather high."787" Including these points does not affect the best-fit values of the model parameters. but indicates that the fit with a CUTOFFPL+BB model is slightly preferable (\2 dofz1.0/104) than the CUTOFFPL+MKL model ,/dofz1.1/104)."," Including these points does not affect the best-fit values of the model parameters, but indicates that the fit with a CUTOFFPL+BB model is slightly preferable $\chi^2_{\rm red}$ /dof=1.0/104) than the CUTOFFPL+MKL model $\chi^2_{\rm red}$ /dof=1.1/104)."788 Further oobservations of this source are probably needed to resolve this issue., Further observations of this source are probably needed to resolve this issue.789 In Fig. 13..," In Fig. \ref{fig:igrbsp},"790 we show the spectrum of aaccumulated over the entire exposure of the observation and fitted with the MKL+CUTOFFPL model., we show the spectrum of accumulated over the entire exposure of the observation and fitted with the MKL+CUTOFFPL model.791 In this figure. we also show for comparison the residuals obtained by fitting the same spectrum with a simple absorbed CUTOFFPL model (I.z101 keV. fixed).," In this figure, we also show for comparison the residuals obtained by fitting the same spectrum with a simple absorbed CUTOFFPL model $E_{\rm cut}$ =11 keV, fixed)."792 The unfolded spectrum of iis shown in Fig. 14.., The unfolded spectrum of is shown in Fig. \ref{fig:igrbunfolded}.793 We searched for pulsations in the power spectra of the oobservation of bby using the same technique described in Sect. 4.1.., We searched for pulsations in the power spectra of the observation of by using the same technique described in Sect. \ref{sec:xteresults}.794 No significant (above 36 level) signal was detected in these data., No significant (above $\sigma$ level) signal was detected in these data.795" We determined an upper limit to the pulsed fraction of 2560,30%.. and for periods in the range 0.3-50 s. 50-100 s. and 0.15-0.3 s. respectively (30 c.1.)."," We determined an upper limit to the pulsed fraction of , and for periods in the range 0.3-50 s, 50-100 s, and 0.15-0.3 s, respectively $\sigma$ c.l.)."796 We have presented the first deep pointed observations of the two prototypical SFXTs. aand In quiescence.," We have presented the first deep pointed observations of the two prototypical SFXTs, and in quiescence."797 The two sources exhibited a very complex timing and spectral variability. and we discuss them separately below.," The two sources exhibited a very complex timing and spectral variability, and we discuss them separately below."798" Here we also carry out a comparison between their quiescent and outburst emission,", Here we also carry out a comparison between their quiescent and outburst emission.799 The two oobservations analyzed1n Sect., The two observations analyzedin Sect.800 4.1 show that the quiescent emission of, \ref{sec:xteresults} show that the quiescent emission of801spectrum (Tribble(1989).. Tao(1995): Maronetal.(2004))).,"spectrum \citet{tri89}, , \citet{tao95}; \citet{mar04}) )."802 The second is the case in which the flow is laminar and the level of conduction inhibition is compared when the field starts Irom initial states of different levels of (angling subject to an imposed temperature difference across an interface., The second is the case in which the flow is laminar and the level of conduction inhibition is compared when the field starts from initial states of different levels of tangling subject to an imposed temperature difference across an interface.803 This second problem is the focus of our preset. paper., This second problem is the focus of our preset paper.804 Using the ASTRODEAR magnetohvdrodynamies code with anisotropic thermal conduction. we investigate (he influence of initial magnetic structure on thermal conduction in an otherwise laminar flow.," Using the ASTROBEAR magnetohydrodynamics code with anisotropic thermal conduction, we investigate the influence of initial magnetic structure on thermal conduction in an otherwise laminar flow."805 The kev questions we address are: (1) does the interface become unstable? (, The key questions we address are: (1) does the interface become unstable? (8062) how fast is the thermal conduction across the interface compared to ihe unmagnetized case?,2) how fast is the thermal conduction across the interface compared to the unmagnetized case?807 We study these questions using different initial magnete configurations imposed on a planar thermal interface to determine how the conduction depends on the amount of field tangling across the interlace., We study these questions using different initial magnetic configurations imposed on a planar thermal interface to determine how the conduction depends on the amount of field tangling across the interface.808 In section 2. we review (he basic equations of MIID with anisotropic thermal conduction.," In section 2, we review the basic equations of MHD with anisotropic thermal conduction."809 In sections 3D and 4 we provide detailed description of the simualtion setup., In sections 3 and 4 we provide detailed description of the simualtion setup.810 In section 5- and 6 we present the simulation resulis and analvses., In section 5 and 6 we present the simulation results and analyses.811 In section 7. we discuss the simulation results in the context of the WBB cooling problem and the cooling flow problem in cores of galaxy. clusters.," In section 7, we discuss the simulation results in the context of the WBB cooling problem and the cooling flow problem in cores of galaxy clusters."812 The appendix provides more detailed information on the testing of the ASTRODEAR code., The appendix provides more detailed information on the testing of the ASTROBEAR code.813 The MILID equations with anisotropic heat conduction that we will solve are givenbv:, The MHD equations with anisotropic heat conduction that we will solve are givenby:814galactic potentialcan decrease the merger time by au order of magnitude and siguificautlv lnerease the ecceutricity of the binary via the “Iwozai mechanism (?)..,galactic potential—can decrease the merger time by an order of magnitude and significantly increase the eccentricity of the binary via the “Kozai mechanism” \citep{blaesleesec02}.815 Iu the following. we will assume that some mechanism Like gas κος is mdeed successful.," In the following, we will assume that some mechanism like gas physics is indeed successful."816 Thus. all relevant pairs cuter the gravitational wave regime aud eventually shed their orbital cnerey iu the form of eyavitational radiation.," Thus, all relevant pairs enter the gravitational wave regime and eventually shed their orbital energy in the form of gravitational radiation."817 The final ten thousand seconds or so of this evolution are thought to be the most enerectic CAV events in the (preseut-«dav). Universe. and will be easily detected by the LISA satellite over a wide range of MDII masses aud redshifts (2)..," The final ten thousand seconds or so of this evolution are thought to be the most energetic GW events in the (present-day) Universe, and will be easily detected by the LISA satellite over a wide range of MBH masses and redshifts \citep{HughesMN2001}."818 Tere. though. we are concerned with the quicsceut evolution that precedes the imfall eveuts.," Here, though, we are concerned with the quiescent quasi-Keplerian evolution that precedes the infall events."819 Now. we will briefly review the cussion of gravitational radiation from binary black holes.," Now, we will briefly review the emission of gravitational radiation from binary black holes."820 First. we will uced the amplitude of the gravitational radiation emitted by a binary at distance D.," First, we will need the amplitude of the gravitational radiation emitted by a binary at distance $D$ ."821" This is even by (?7) where Mo=[MyARAL,|MS).O7 is the ~chirp mass” of the svsteii aud D, is the proper rest-frame period of the binary."," This is given by \citep{Peters63,Thorne}822 where ${\cal M}=[M_1 M_2 (M_1 + M_2)^{-1/3}]^{3/5}$ is the “chirp mass” of the system, and $P_p$ is the proper rest-frame period of the binary."823" The observed frequency of the radiation iu the ath harmonic will be f=»f,/(l|2). where f,=1/P,."," The observed frequency of the radiation in the $n$ th harmonic will be $f=nf_p/(1+z)$, where $f_p\equiv 1/P_p$."824 Ouly terms with »=2 are nonzero. aud v=2 for the circular orbits to which we will restrict ourselves hereafter (??)..," Only terms with $n\ge2$ are nonzero, and $n=2$ for the circular orbits to which we will restrict ourselves hereafter \citep{quinlan96,quinlanhernquist97}."825 However. if the binaries are driven together by the presence of a third MDIT (?).. the eccentricity may be significant aud populate highertrequency harmonics.," However, if the binaries are driven together by the presence of a third MBH \citep{blaesleesec02}, the eccentricity may be significant and populate higher-frequency harmonics."826 This would oulv serve to increase the amplitude of gravitational radiation aud populate the ligher harmouics (bv au factor)., This would only serve to increase the amplitude of gravitational radiation and populate the higher harmonics (by an eccentricity-dependent factor).827 This amplification would offset. or exceed. the loss resulting frou expulsion of the third body which is assmuce to coalescence in our model.," This amplification would offset, or exceed, the loss resulting from expulsion of the third body which is assumed to coalescence in our model."828 Equatiou 21 differs by a factor of y‘3/7 [from that eiven in ?.., Equation \ref{eq:strain} differs by a factor of $\sqrt{3/4}$ from that given in \citetalias{Rajagopal95}.829 Instead of the “characteristic strain” from ?.. which includes this factor to recover the signal-to-noise for an carth-bound detector. we just use the augle-averaged mican-square strain(ie... eq. [2 Uf ," Instead of the “characteristic strain” from \citet{Thorne}, which includes this factor to recover the signal-to-noise for an earth-bound detector, we just use the angle-averaged mean-square strain, eq. \ref{eq:strain}] ]"830is 1/2 times the maxim rms strain)., is $1/\sqrt{2}$ times the maximum rms strain).831 We point out for ease of interpretation of this aud later formmiaec that (GMc7) has units of time., We point out for ease of interpretation of this and later formulae that $(GM/c^3)$ has units of time.832 In a cosmological setting. we replace D with στρΕμ.," In a cosmological setting, we replace $D$ with $c a_0 H_0833r(z)/H_0$."834 We will also need the characteristic timescale of the cussion. defined im equation (9)).," We will also need the characteristic timescale of the emission, defined in equation \ref{eq:taugwdef}) )."835" Usine the Iepler- formmla.. fra?PM=(22).oολAM») and the formmla for enerev loss from gravitational waves. the eyavitational wave timescale is (e.g.7?) where we have used f,=fl|:)/n aud sot a=2 for circular orbits."," Using the Kepler formula, $f^2 a^3 = (2\pi)^{-2} G(M_1+M_2)$ and the formula for energy loss from gravitational waves, the gravitational wave timescale is \citep[\eg,][]{Peters63,ShapiroTeukolsky}836 where we have used $f_p=f(1+z)/n$ and set $n=2$ for circular orbits."837 Now we can combine all of theingredients iuto au observable strain power spectrum., Now we can combine all of theingredients into an observable strain power spectrum.838 We have, We have839counter windows.,counter windows.840" The sources of a different spectral shape than the Crab's, then, may well have greater scatter."," The sources of a different spectral shape than the Crab's, then, may well have greater scatter."841" Second, the Cygnus region harbours the very bright and variable source Cyg Χ-Ι, which might contribute to additional systematic error in the ASM data."," Second, the Cygnus region harbours the very bright and variable source Cyg X-1, which might contribute to additional systematic error in the ASM data."842 The results based on the data revolution 567-569 have been combined as there were no differences within the error bars., The results based on the data revolution 567-569 have been combined as there were no differences within the error bars.843" In Fig.2 we show the IBIS/ISGRI spectrum of revolution 567-569, which shows that Cygnus X-3 is indeed in the low hard state again."," In \ref{fig:lowhard} we show the IBIS/ISGRI spectrum of revolution 567-569, which shows that Cygnus X-3 is indeed in the low hard state again."844 The spectrum can be described by a single power-law model with photon index Γ=3.94+0.03., The spectrum can be described by a single power-law model with photon index $\Gamma = 3.94 \pm 0.03$.845 It is apparent that the transition from the ultrasoft state to the low hard state is accompanied by a steepening of the high-energy component., It is apparent that the transition from the ultrasoft state to the low hard state is accompanied by a steepening of the high-energy component.846" At the same time, the luminosity of this component has increased by a factor of ~14."," At the same time, the luminosity of this component has increased by a factor of $\sim 14$."847 A radio observation within the transition phase on 2007-06-01 (Trushkin et al., A radio observation within the transition phase on 2007-06-01 (Trushkin et al.848" 20078) showed strong emission (>1 Jy), which indicates increased jet activity."," 2007a) showed strong emission $>8491$ Jy), which indicates increased jet activity."850 The behaviour during this transition from the ultrasoft to low hard states is comparable to what was reported in January 2007 (Trushkin et al., The behaviour during this transition from the ultrasoft to low hard states is comparable to what was reported in January 2007 (Trushkin et al.851 2007b)., 2007b).852" The ultrasoft state is apparent during spacecraft revolutions 562 and 563, with an energy output of Efg=0.05keVcm?s! in the 20 — 60 keV energy band."," The ultrasoft state is apparent during spacecraft revolutions 562 and 563, with an energy output of $E f_E = 0.05 \rm \, keV \, cm^{-2} \,853s^{-1}$ in the 20 – 60 keV energy band."854" In revolution 564, we observed the onset of the transition, with a steeper spectral slope in the hard X-rays but still comparably low luminosity."," In revolution 564, we observed the onset of the transition, with a steeper spectral slope in the hard X-rays but still comparably low luminosity."855 Figure 3 shows the evolution of the spectral slope and of the luminosity in the 20-60 keV energy band with time., Figure \ref{fig:evolution} shows the evolution of the spectral slope and of the luminosity in the 20–60 keV energy band with time.856" In revolution 567, the low hard state is reached and the ISGRI data can be modelled by a simple power law with Il=3.94+ 0.03."," In revolution 567, the low hard state is reached and the ISGRI data can be modelled by a simple power law with $\Gamma = 3.94 \pm857 0.03$ ."858" The low hard state has been measured in this energy range before, e.g. by RXTE/HEXTE (Choudhury Rao 2002; Choudhury et al."," The low hard state has been measured in this energy range before, e.g. by /HEXTE (Choudhury Rao 2002; Choudhury et al."859 2002)., 2002).860" The data with higher significance allow them to apply a more complex model, ie. Comptonization of seed photons from a thermal multi- accretion disk by a thermal Comptonizing plasma cloud (CompST; Sunyaev&Titarchuk 1980)) with electron temperature kT,=4.9+0.1keV and a single power law with I=2.01x0.04."," The data with higher significance allow them to apply a more complex model, i.e. Comptonization of seed photons from a thermal multi-coloured accretion disk by a thermal Comptonizing plasma cloud (CompST; \cite{CompST}) ) with electron temperature $kT_e861 = 4.9 \pm 0.1 \rm \, keV$ and a single power law with $\Gamma = 2.01 \pm 0.04$."862" Although the ISGRI data are described well by a simple power law with Γ= 3.9, we applied the more complex model for comparison reasons, which results in kT,=5.2+4.0keV and I=3.4€ 1.0."," Although the ISGRI data are described well by a simple power law with $\Gamma = 3.9$ we applied the more complex model for comparison reasons, which results in $kT_e = 5.2 \pm 4.0 \rm \, keV$ and $\Gamma = 3.4 \pm 1.0$ ."863" When freezing the power law to the value reported by Choudhury Rao, the electron temperature becomes KT,=8.6+2.1keV."," When freezing the power law to the value reported by Choudhury Rao, the electron temperature becomes $kT_e = 8.6 \pm 2.1 \rm \, keV$."864" Thus, to explain the observed ISGRI spectrum in terms of a Comptonization component plus a power law, a stronger Comptonization component is observed here than by RXTE/HEXTE in May 1998."," Thus, to explain the observed ISGRI spectrum in terms of a Comptonization component plus a power law, a stronger Comptonization component is observed here than by /HEXTE in May 1998."865" It has to be kept in mind, though,thatthe ISGRI data of the low hard state alone"," It has to be kept in mind, though,thatthe ISGRI data of the low hard state alone"866The Magellanic Clouds. which are the largest nearby ealaxies. provide our most detailed view of the extragalactic universe.,"The Magellanic Clouds, which are the largest nearby galaxies, provide our most detailed view of the extragalactic universe."867 Although their proximity is eoncrally an advantage. their laree augular exteut on the sky has hampered global studies.," Although their proximity is generally an advantage, their large angular extent on the sky has hampered global studies."868 Historically. stellar catalogs of the Clouds have relied on photographic data (sce Hatzidinitriouctal.(1989). and Tnwinetal.(1990) for some of the most recent examples).," Historically, stellar catalogs of the Clouds have relied on photographic data (see \cite{hatz89} and \cite{idk90} for some of the most recent examples)."869 Within the last decade. several large-scale digital survevs of the Magellanic Clouds have been undertaken.," Within the last decade, several large-scale digital surveys of the Magellanic Clouds have been undertaken."870 In the optical bands. the principal ones are the microleusing surveys (MACTIO. Alcocketal.(1997): OGLE. Udalskietal.(1998): aud EROS. Palanque-Delabrouilleetal. (1998))). an eniissiou lue survey (Sunithetal. (2000))). a bright star survev (Massey (2001))). the red optical channel (£j) of the infrared DENIS survey (Epchteiuetal. (1997))). and our Magellanic Clouds Photometric Survey. hereafter: ALCPS (Zaritskyctal. (1997))).," In the optical bands, the principal ones are the microlensing surveys (MACHO, \cite{alcock97}; OGLE, \cite{udalski98}; and EROS, \cite{p98}) ), an emission line survey \cite{smith00}) ), a bright star survey \cite{massey01}) ), the red optical channel $I$ ) of the infrared DENIS survey \cite{ep97}) ), and our Magellanic Clouds Photometric Survey, hereafter MCPS \cite{zht97}) )."871 We present the stellar photometric data in catalog form from the ALICPS for the cutive Small Magellanic Cloud (SAIC) survey region (rouelly 1.57«E. where the longer direction is north-south).," We present the stellar photometric data in catalog form from the MCPS for the entire Small Magellanic Cloud (SMC) survey region (roughly $4.5^\circ \times 4^\circ$, where the longer direction is north-south)."872 The principal advantages of these data in comparison to the surveys listed above are that our data are either deeper. cover a wider area. or include a larger number of filters (the inclusion of ( is particularly iuportant for studies of dust aud young stellar populations).," The principal advantages of these data in comparison to the surveys listed above are that our data are either deeper, cover a wider area, or include a larger number of filters (the inclusion of $U$ is particularly important for studies of dust and young stellar populations)."873 However. as we show iu our comparison to these other catalogs. each of the other surveys has its complementary streneths aud we incorporate data frou several of them to augment the MCDPS catalog.," However, as we show in our comparison to these other catalogs, each of the other surveys has its complementary strengths and we incorporate data from several of them to augment the MCPS catalog."874 Iu addition to providing the catalog. we construct aud analyze extinction maps of the SMC.," In addition to providing the catalog, we construct and analyze extinction maps of the SMC."875 As we demonstrated or a portion of the LAIC (Zaxitsky (1999))). the extinction xoperties in the Clouds are not only spatially variable. mit depend ou stellar population.," As we demonstrated for a portion of the LMC \cite{z99}) ), the extinction properties in the Clouds are not only spatially variable, but depend on stellar population."876 Therefore. for miu scientific purposes the catalog aloue is iusufficieut. one uust correct the observed imnagnitudes and colors for a complex extiuctiou pattern.," Therefore, for many scientific purposes the catalog alone is insufficient, one must correct the observed magnitudes and colors for a complex extinction pattern."877 We describe the AICPS in 822. discuss detailed. comparisons with previous daa to assess he quality of the catalog in 8323. use the photometry to eonerate extinction maps of the SAIC for two different stellar populations iu Stl. and present the final catalog πι 85h.," We describe the MCPS in 2, discuss detailed comparisons with previous data to assess the quality of the catalog in 3, use the photometry to generate extinction maps of the SMC for two different stellar populations in 4, and present the final catalog in 5."878 The data come from the ongoing Alagellanic Cloud Photometric Survey (Zaritskyetal. (1997)))., The data come from the ongoing Magellanic Cloud Photometric Survey \cite{zht97}) ).879 Usine the, Using the880 Inrecent vears. an enormous amount of exciting cosmological data have appeared. accompanied by theoretical inferences about early galaxy formation and the first massive stars.,"In recent years, an enormous amount of exciting cosmological data have appeared, accompanied by theoretical inferences about early galaxy formation and the first massive stars."881 May of (hese inferences were reactions lo first-vear (WAIAP-1) results (IXogut. et 22003: Spergel 2003) from the (WNLAD)., Many of these inferences were reactions to first-year (WMAP-1) results (Kogut et 2003; Spergel 2003) from the (WMAP).882 WAIAP-1 inferred. a high optical depth to the cosmic microwave backeround (CAIB) and suggested early reionization of (he intergalactic medium (IGM)., WMAP-1 inferred a high optical depth to the cosmic microwave background (CMB) and suggested early reionization of the intergalactic medium (IGM).883 Other conclusions came from simplified moclels for the stellar initial mass function (AIF). atomic/molecular physics. radiative processes. and prescriptions for star formation rates aud escape of photoionizing radiation [rom protogalaxies.," Other conclusions came from simplified models for the stellar initial mass function (IMF), atomic/molecular physics, radiative processes, and prescriptions for star formation rates and escape of photoionizing radiation from protogalaxies."884 The CAIB optical depth and other cosmological parameters have been refined significantly in (he recent WMAP-5 data (Llinshaw 2008)., The CMB optical depth and other cosmological parameters have been refined significantly in the recent WMAP-5 data (Hinshaw 2008).885 In this paper. we use these new measurements lo constrain the efficiency of first-lisht ionizing sources.," In this paper, we use these new measurements to constrain the efficiency of first-light ionizing sources."886" We focus on the reionization epoch. defined as the redshilt 2, when the IGM becomes nearly fully ionized over most of its volume (Gnedin 2000. 2004)."," We focus on the reionization epoch, defined as the redshift $z_r$ when the IGM becomes nearly fully ionized over most of its volume (Gnedin 2000, 2004)."887 Our knowledge about reionization comes primarily [from (ree (vpes of observations: hydrogen aabsorption in the IGM. hieh-: Lya-emittine galaxies. and CAIB optical depth.," Our knowledge about reionization comes primarily from three types of observations: hydrogen absorption in the IGM, $z$ -emitting galaxies, and CMB optical depth."888" Optical spectroscopic studies of the ""Gunn-Peterson (Lva)) absorption toward high-redshift quasars and galaxies imply that IL I reionization occurred not far bevond στον~6 (Becker 2001: Fan 2002. 2006)."," Optical spectroscopic studies of the “Gunn-Peterson"" ) absorption toward high-redshift quasars and galaxies imply that H I reionization occurred not far beyond $z_{\rm GP} \sim 6$ (Becker 2001; Fan 2002, 2006)."889 Ultraviolet spectra suggest that Ile II reionization occurred at z3 (Ixriss 2001: Shull 2004: Zheng 2004)., Ultraviolet spectra suggest that He II reionization occurred at $z \sim 3$ (Kriss 2001; Shull 2004; Zheng 2004).890" The detection of high-redshift -emitting) galaxies (Ili Cowie 2006) suggests a somewhat higher redshift. 2,>6.5."," The detection of high-redshift -emitting) galaxies (Hu Cowie 2006) suggests a somewhat higher redshift, $z_r \geq 6.5$."891 Data from WMAP. alter three vears (Spergel 2007) and five vears (Hinshaw 2008). suggest (hat reionization might occur at z£10. with sizeable uncertainties in measuring ancl modeling the CAIB optical depth.," Data from WMAP, after three years (Spergel 2007) and five years (Hinshaw 2008), suggest that reionization might occur at $z \approx 10$, with sizeable uncertainties in measuring and modeling the CMB optical depth."892" There appears to be a cliscrepancy between (he (wo epochs. cz 6T and z,zLO."," There appears to be a discrepancy between the two epochs, $\approx$ 6–7 and $z_r \approx 10$."893" However. the WMAP and aabsorption results are not necessarily inconsistent. since (μον probe small amounts of ionized and neutral gas. respectively,"," However, the WMAP and absorption results are not necessarily inconsistent, since they probe small amounts of ionized and neutral gas, respectively."894" Both the HE I absorbers and ionized filaments in the ""cosmic web"" (Cen Ostriker 1999) are hiehlv structured al redshifts z«10 and affect the optical depths inLvo."," Both the H I absorbers and ionized filaments in the “cosmic web"" (Cen Ostriker 1999) are highly structured at redshifts $z < 10$ and affect the optical depths in."895". In order to effectively absorb all the radiation al z£z6 requires a volume-averaged neutral fraction of just ry,z4xLO+ (Fan 2006).", In order to effectively absorb all the radiation at $z \approx 6$ requires a volume-averaged neutral fraction of just $x_{\rm HI} \approx 4 \times 10^{-4}$ (Fan 2006).896 Simulations of the reionization process (Gnedin 2004: Gnedin Fan 2006) show that the transition from neutral to ionized is extended in time between 2=5—10., Simulations of the reionization process (Gnedin 2004; Gnedin Fan 2006) show that the transition from neutral to ionized is extended in time between $z = 5-10$.897 The first stage (pre-overlap) involves the development. and expansion of the first isolated ionizing sources., The first stage (pre-overlap) involves the development and expansion of the first isolated ionizing sources.898 The second stage marks (he overlap of the ionization fronts and (he disappearance of the last vestiges of low-density neutral gas., The second stage marks the overlap of the ionization fronts and the disappearance of the last vestiges of low-density neutral gas.899 Finally. in the post-overlap stage. the remaining," Finally, in the post-overlap stage, the remaining"900standard deglitching and crosstalk correction were not used due to poor results in these fields. and custom procedures were written for drift removal.,"standard deglitching and crosstalk correction were not used due to poor results in these fields, and custom procedures were written for drift removal."901 SPIRE data processing required less deviation. from the standard processing methods (2).. with both standard deghitching and drift removal producing acceptable results.," SPIRE data processing required less deviation from the standard processing methods \citep{spire}, with both standard deglitching and drift removal producing acceptable results."902 In both cases. the ROMAGAL Generalised Least Squares algorithm (?) was used to produce the final maps.," In both cases, the ROMAGAL Generalised Least Squares algorithm \citep{traficante10} was used to produce the final maps."903 A more thorough discussion of the entire data reduction process can be found in ?.., A more thorough discussion of the entire data reduction process can be found in \citet{traficante10}.904" Each of the S06 sources in these regions was studied at 6 wavebands. Le. GLIMPSE 8,;m (??) and the five Hi-GAL wavelengths (70. 160. 250. 350 and um)."," Each of the S06 sources in these regions was studied at 6 wavebands, i.e. GLIMPSE $\mu$ m \citep{glimpse, spitzer} and the five Hi-GAL wavelengths (70, 160, 250, 350 and $\mu$ m)."905 If the source showed extinction at MIR but emission in the FIR it was identified as an IRDC., If the source showed extinction at MIR but emission in the FIR it was identified as an IRDC.906 Within the region. approximately 330 (PF09) IRDCs were found. each containing at least one infrared dark core.," Within the region, approximately 330 (PF09) IRDCs were found, each containing at least one infrared dark core."907 Our aim is to model those cores without embedded protostars., Our aim is to model those cores without embedded protostars.908 We focus on three objects. GO030.50+00.95. G031.03400.26 and G031.03-00.76. which meet the criteria for IRDCs.," We focus on three objects, G030.50+00.95, G031.03+00.26 and G031.03+00.76, which meet the criteria for IRDCs."909" All three contain two cores each. designated ""Core A and ""Core B' in order of Right Ascension within each cloud."," All three contain two cores each, designated `Core A' and `Core B' in order of Right Ascension within each cloud."910 Positions and physical properties of all six cores can be found in Table 1.. and images are displayed in Figs. 1-—3..," Positions and physical properties of all six cores can be found in Table \ref{coreprop}, , and images are displayed in Figs. \ref{model3095}- \ref{model3176}."911 — this cloud. Core B is the more extended and brighter of the two cores and located to the south ofCore A (see Fig. 1)).," In this cloud, Core B is the more extended and brighter of the two cores and located to the south ofCore A (see Fig. \ref{model3095}) )."912 — the 250m and 350m wavebands. it is possible to see that Core B can be separated into two cores.," In the $\mu$ m and $\mu$ m wavebands, it is possible to see that Core B can be separated into two cores."913 However. as the individual peaks cannot be seen at either. 160m or 500m. the two are grouped together and modelled as a single core.," However, as the individual peaks cannot be seen at either $\mu$ m or $\mu$ m, the two are grouped together and modelled as a single core."914 A distance of kkpe for the entire cloud is assumed (?).., A distance of kpc for the entire cloud is assumed \citep{simon06b}.915 o previously calculated masses are available for these cores., No previously calculated masses are available for these cores.916 Results calculated from our data are presented in Table 1.., Results calculated from our data are presented in Table \ref{coreprop}.917 S06 find three cores in this cloud., S06 find three cores in this cloud.918 However. only two can be seen as significant emission sources in the FIR and so here the third is ignored.," However, only two can be seen as significant emission sources in the FIR and so here the third is ignored."919 Core B is the more extended and brighter of the two and is located to the north of Core A (see Fig. 2))., Core B is the more extended and brighter of the two and is located to the north of Core A (see Fig. \ref{model3126}) ).920 Emission from Core B can be clearly seen at wavelengths as short as zm. whereas emission from Core A is not seen shortward of jum. We take a distance of kkpe for both cores (?) although we note that there are two emission line components among the line of sight. with the second emission line giving a distance of kkpe.," Emission from Core B can be clearly seen at wavelengths as short as $\mu$ m, whereas emission from Core A is not seen shortward of $\mu$ m. We take a distance of kpc for both cores \citep{teyssier02} although we note that there are two emission line components among the line of sight, with the second emission line giving a distance of kpc."921 Therefore there is some uncertainty in this distance., Therefore there is some uncertainty in this distance.922 ? use 850m data and find masses of 420 and MM... respectively. for these cores.," \citet{parsons09} use $\mu$ m data and find masses of 420 and $_{\odot}$ , respectively, for these cores."923 However. these authors used adifferent size of aperture.," However, these authors used adifferent size of aperture."924 We make a direct comparison with, We make a direct comparison with925"of the spherical molecular core of the cloud, which is surrounded by a shell of atomic gas whose density is lower than that of the molecular gas by a factor mol, which KMTO09 show is typically ~10.","of the spherical molecular core of the cloud, which is surrounded by a shell of atomic gas whose density is lower than that of the molecular gas by a factor $\phi_{\rm mol}$, which \citetalias{krumholz09a} show is typically $\simeq 10$."926" The covering fraction of the molecular sphere is Obviously a spherical ball is a great oversimplification of the complex geometries of atomic-molecular complexes, but cy, is a useful general indicator of the fraction of the area that is likely to be covered by molecular material."," The covering fraction of the molecular sphere is Obviously a spherical ball is a great oversimplification of the complex geometries of atomic-molecular complexes, but $c_{\rm H_2}$ is a useful general indicator of the fraction of the area that is likely to be covered by molecular material."927For a given metallicity Z' it is trivial to numerically invert equation (2)) to calculate the total cold gas column density Νο for which the molecular covering fraction reaches a particular value ,For a given metallicity $Z'$ it is trivial to numerically invert equation \ref{coveringfrac}) ) to calculate the total cold gas column density $N_{\rm c}$ for which the molecular covering fraction reaches a particular value $c_{\rm H_2}$ .928"The mean atomic column density is then N(H1)=cy,.(1— fu,)Nec.", The mean atomic column density is then $N(\hi) = (1-f_{\rm H_2}) N_{\rm c}$ .929" This defines a locus of points in the N(H1), Z'-plane corresponding to the specified cy,."," This defines a locus of points in the $N(\hi), Z'$ -plane corresponding to the specified $c_{\rm H_2}$."930" The maximum ccolumn density corresponds to the limit cy,—1, because this corresponds to an infinite slab illuminated by the external radiation field."," The maximum column density corresponds to the limit $c_{\rm H_2} \rightarrow 1$, because this corresponds to an infinite slab illuminated by the external radiation field."931 At this point we must mention two important caveats., At this point we must mention two important caveats.932" One is that we assume that the atomic gas in DLAs is in two-phase equilibrium, which may not be true for all of them."," One is that we assume that the atomic gas in DLAs is in two-phase equilibrium, which may not be true for all of them."933 The other is that this method allows us to constrain only the cold ccolumn density., The other is that this method allows us to constrain only the cold column density.934 In principle much larger warm gas column densities are possible 9))., In principle much larger warm gas column densities are possible \ref{twophase}) ).935" In Figure 1, we plot our derived values N(H1) versus Z’ for molecular covering fractions from cg,= We also show lines of constant E(B— V), computed using a Draine(2003) Ry=3.1 extinction curve scaled by metallicity, giving E(B—V)/N(H1)=1.65x10-222’ cm? and =5.32x107722’ mag cm?."," In Figure \ref{dlamol_nhz}, we plot our derived values $N(\hi)$ versus $Z'$ for molecular covering fractions from $c_{\rm H_2}=0.01-1$ We also show lines of constant $E(B-V)$ , computed using a \citet{draine03a} $R_V=3.1$ extinction curve scaled by metallicity, giving $E(B-V)/N(\hi) = 1.65\times 10^{-22} Z'$ $^2$ and $A_V/N(\hi)=5.32\times 10^{-22} Z'$ mag $^2$."936" We compare toAy/N(H1) observed QSO- and GRB-from Herbert-Fortetal.(2006),, Prochaskaetal.(2007, 2009),, Kaplan et ((2009, in preparation), and Dessauges-Zavadsky et ((2009, in"," We compare to observed QSO- and GRB-DLAsfrom \citet{herbert-fort06a}, \citet{prochaska07a, prochaska09a}, Kaplan et (2009, in preparation), and Dessauges-Zavadsky et (2009, in preparation)."937" For the Dessauges-Zavadsky et ssamplepreparation). we derive metallicities from zinc abundance: logZ’=[Zn/H]— where +12=4.63 log(Zn/H)(Lodderslog(Zn/H)o, For all other"," For the Dessauges-Zavadsky et sample we derive metallicities from zinc abundance: $\log Z' = [{\rm Zn}/{\rm H}] \equiv \log({\rm Zn}/{\rm H}) - \log ({\rm Zn}/{\rm H})_{\odot}$ , where $\log({\rm Zn}/{\rm H})_{\odot} + 12 = 4.63$ \citep{lodders03a}."938 datalog(Zn/H)c we use the metallicity reported by 2003)..the authors., For all other data we use the metallicity reported by the authors.939" To avoid possible issues arising from either ionization correction or metallicity evolution with redshift, we exclude DLAs with logN(H1)«20 and redshift z«1.7."," To avoid possible issues arising from either ionization correction or metallicity evolution with redshift, we exclude DLAs with $\log N(\hi)<20$ and redshift $z<1.7$."940" As the Figure shows, the zone at high N(H1) and Z' where no DLAs lie (except that associated with GRBO080607, which we discuss below), corresponds well to the predicted zone of exclusion."," As the Figure shows, the zone at high $N(\hi)$ and $Z'$ where no DLAs lie (except that associated with GRB080607, which we discuss below), corresponds well to the predicted zone of exclusion."941" The molecular covering fraction declines sharply away from the 1 line, so all DLAs but GRB080607 lie below 0.06."," The molecular covering fraction declines sharply away from the $c_{\rm H_2} = 1$ line, so all DLAs but GRB080607 lie below $c_{\rm H_2}=0.06$ ."942" This is consistent with the results of Zwaancy,&Prochaska (2006),, who conclude that detection of true molecular clouds in DLAs is unlikely because the molecular material has a small covering fraction."," This is consistent with the results of \citet{zwaan06a}, who conclude that detection of true molecular clouds in DLAs is unlikely because the molecular material has a small covering fraction."943" Trace amounts of molecular hydrogen have been discovered in some DLAs (e.g.Ledouxetal.2003;Noterdaemeetal. 2008),, but these low molecular columns almost certainly correspond to Hs spatially mixed with cold atomic gas, rather than true molecular clouds."," Trace amounts of molecular hydrogen have been discovered in some DLAs \citep[e.g.][]{ledoux03a, noterdaeme08a}, but these low molecular columns almost certainly correspond to $_2$ spatially mixed with cold atomic gas, rather than true molecular clouds."944" The KMT formalism approximates the atomic-molecular transition as sharp, so it does not apply to these systems."," The KMT formalism approximates the atomic-molecular transition as sharp, so it does not apply to these systems."945 We defer discussion of them to future work., We defer discussion of them to future work.946" Also note that the observed distribution falls off sharply at logZ’=0, and at logN(H1)=22 independent of Z’. Molecule formation cannotexplainthese features."," Also note that the observed distribution falls off sharply at $\log Z' \ga 0$, and at $\log N(\hi) \ga 22$ independent of $Z'$ Molecule formation cannotexplainthese features."947" The DLA associated with GRB080607 (Prochaska 2009),, theonly DLA inside the zone of exclusion, is also the only DLA to show significant columns ofH5 and CO."," The DLA associated with GRB080607 \citep{prochaska09a}, , theonly DLA inside the zone of exclusion, is also the only DLA to show significant columns of$_2$ and CO."948 We plot this detection at two metallicities derived in different ways., We plot this detection at two metallicities derived in different ways.949 The logZ’=—0.2 point corresponds to, The $\log Z'=-0.2$ point corresponds to950which give rise to subhaloes are more concentrated than those of the larger haloes they fall into.,which give rise to subhaloes are more concentrated than those of the larger haloes they fall into.951 In addition. we show in the next section that the radial distribution of subhaloes is less concentrated than that of the mass.," In addition, we show in the next section that the radial distribution of subhaloes is less concentrated than that of the mass."952 Both these effects should. reduce the dillerence between the mass assigned to an isolated. halo and that assigned to the subhalo it. turns into., Both these effects should reduce the difference between the mass assigned to an isolated halo and that assigned to the subhalo it turns into.953 On the other hand. dynamical processes strip material from a halo once it is incorporated into a larger system. thereby reducing its mass.," On the other hand, dynamical processes strip material from a halo once it is incorporated into a larger system, thereby reducing its mass."954 As we demonstrate in Section 5. most subhaloes fell into their host relatively recently and the amount of stripping is typically quite modest.," As we demonstrate in Section 5, most subhaloes fell into their host relatively recently and the amount of stripping is typically quite modest."955 The combined ellect of all these factors is that once subhalo masses are doubled. as above. the number of subhaloes per unit miss within a halo is verv similar to the number of small haloes per unit mass in the surrounding universe ancl thus in the material [rom which the main halo formed.," The combined effect of all these factors is that once subhalo masses are doubled, as above, the number of subhaloes per unit mass within a halo is very similar to the number of small haloes per unit mass in the surrounding universe and thus in the material from which the main halo formed."956 —ow are subhaloes distributed: within their parent halo?, How are subhaloes distributed within their parent halo?957 uluperficiallv. this appears closely related to the distribution of galaxies within clusters. but in fact this relation is complicated because subhalo masses are much more strongly alfected by. tidal stripping than are the luminosities of the ealaxies they contain.," Superficially, this appears closely related to the distribution of galaxies within clusters, but in fact this relation is complicated because subhalo masses are much more strongly affected by tidal stripping than are the luminosities of the galaxies they contain."958 As a result the effective total mass-to-light> ratio of cluster ogalaxies is a stronglye increasingIn function of clustercentric radius (see Fig., As a result the effective total mass-to-light ratio of cluster galaxies is a strongly increasing function of clustercentric radius (see Fig.959 12 of SWThs)., 12 of SWTK).960 lt ds uso interesting to ask whether the radial distribution of subhaloes depends on subhalo mass or on the mass of the parent halo., It is also interesting to ask whether the radial distribution of subhaloes depends on subhalo mass or on the mass of the parent halo.961 We address the latter dependence using haloes from our CGIE2 and. cluster simulations split into the three mass ranges already analysed. in Section 4.2., We address the latter dependence using haloes from our GIF2 and cluster simulations split into the three mass ranges already analysed in Section 4.2.962 For cach mass range we compute the mean fraction »» number of all subhaloes within req) that [ie within normalized: radius r£rogg., For each mass range we compute the mean fraction by number of all subhaloes within $r_{200}$ that lie within normalized radius $r/r_{200}$.963 Dhese subhalo number density xofiles are shown in the upper left-hand. panel of Fie., These subhalo number density profiles are shown in the upper left-hand panel of Fig.964 10 and are compared with a similarly defined. profile for the otal mass., 10 and are compared with a similarly defined profile for the total mass.965 All data are shown for z=0 and for subhaloes with maaMyygs20.001 only.," All data are shown for $z=0$ and for subhaloes with $m_{\rm966sub}/M_{\rm halo}> 0.001$ only."967 We can then get comparable and reliable results for all three halo mass ranges., We can then get comparable and reliable results for all three halo mass ranges.968 Lt is clear hat the racial distribution of subhaloes is substantially less concentrated than that of the mass as a whole., It is clear that the radial distribution of subhaloes is substantially less concentrated than that of the mass as a whole.969" There is no significant. dependence detected. on parent. halo mass over he one order of magnitudeo rangeὃν tested in this panel. but a weak dependence does appear when we compare with our ""Milky: Was’ simulation GASH (see below)."," There is no significant dependence detected on parent halo mass over the one order of magnitude range tested in this panel, but a weak dependence does appear when we compare with our `Milky Way' simulation GA3n (see below)."970 We acldress the issue of possible dependences on subhalo mass using our cluster resimulations together with the haloes in the most massive bin of our GIE2 simulation (for a total of 15 systems)., We address the issue of possible dependences on subhalo mass using our cluster resimulations together with the haloes in the most massive bin of our GIF2 simulation (for a total of 15 systems).971 In the upper right-hand. panel of Fig., In the upper right-hand panel of Fig.972 LO we show radial number fraction plots for subhalo populations limited above LO* and 10.1 of the parent halo mass., 10 we show radial number fraction plots for subhalo populations limited above $10^{-3}$ and $10^{-4}$ of the parent halo mass.973 There appears to be a slight tendeney for the more massive haloes to be more centrally concentrated. but the cllect is small and it is unclear i£ it is significant given the relatively small number of parent haloes in our sample.," There appears to be a slight tendency for the more massive haloes to be more centrally concentrated, but the effect is small and it is unclear if it is significant given the relatively small number of parent haloes in our sample."974 For these same 15 clusters. the upper right-hand panel of Fig.," For these same 15 clusters, the upper right-hand panel of Fig."975 10 also shows the cumulative racial profile of subhaloes for which Vis is greater than 10 per cent of the parent halo’s value. of Vooo., 10 also shows the cumulative radial profile of subhaloes for which $V_{\rm max}$ is greater than 10 per cent of the parent halo's value of $V_{200}$.976 I ds interesting that. this population appears to be significantly more concentrater than populations defined in these same haloes above a mass threshold., It is interesting that this population appears to be significantly more concentrated than populations defined in these same haloes above a mass threshold.977 “Phis presumably results from a combination of two effects., This presumably results from a combination of two effects.978 A subhalo of given density structure is assignee smaller ancl smaller masses but larecly unchanging Visas values as it gets closer to the centre of its parent halo., A subhalo of given density structure is assigned smaller and smaller masses but largely unchanging $V_{\rm max}$ values as it gets closer to the centre of its parent halo.979 In addition. subhaloes near the centre of their parent tend to be more heavily allected by tidal μαstripping than more clistant objects.," In addition, subhaloes near the centre of their parent tend to be more heavily affected by tidal stripping than more distant objects."980 As demonstrated in Section 5.4. such tidal stripping allects the masses of subhaloes more strongly than their maximum circular velocities (Chigna ct al.," As demonstrated in Section 5.4, such tidal stripping affects the masses of subhaloes more strongly than their maximum circular velocities (Ghigna et al."981 2000: Hayashi et al., 2000; Hayashi et al.982 2003: WKravisov et al., 2003; Kravtsov et al.983 20045)., 2004b).984 The lower panels of Fig., The lower panels of Fig.985 10 use our Milkv Waa’ simulations to extend. these results to. parent haloes of lower mass and to test further for resolution ellects., 10 use our `Milky Way' simulations to extend these results to parent haloes of lower mass and to test further for resolution effects.986 The dashed. ancl solid. curves compare the cumulative profiles for subhaloes with mass greater than 5.78«105.1M. in GA? and GASn., The dashed and solid curves compare the cumulative profiles for subhaloes with mass greater than $5.78 \times 10^7h^{-1}{\rm M_\odot}$ in GA2 and GA3n.987 This mass corresponds to 30 particles in GA? and is μην0000.," This mass corresponds to 30 particles in GA2 and is $M_{\rm988halo}/40000$."989 “Phe two profiles agree extremely well. suggesting that resolution is not seriously elfecting our subhalo cistributions.," The two profiles agree extremely well, suggesting that resolution is not seriously effecting our subhalo distributions."990 Recucing the lower limit on subhalo particle number still further does lead to noticeable ellects. as we show in the lower right-hand panel of Fig.," Reducing the lower limit on subhalo particle number still further does lead to noticeable effects, as we show in the lower right-hand panel of Fig."991 10., 10.992 l]lere the comparison is repeated for the subhalo mass range Corresponding to 10 to 30 particles in C:A2., Here the comparison is repeated for the subhalo mass range corresponding to 10 to 30 particles in GA2.993 The abundance of subhaloes is significantly depressed. in the lower resolution simulation. particularly in the inner regions.," The abundance of subhaloes is significantly depressed in the lower resolution simulation, particularly in the inner regions."994 Near the resolution limit of a simulation subhaloes begin to be lost and they disappear preferentially in the inner regions of haloes., Near the resolution limit of a simulation subhaloes begin to be lost and they disappear preferentially in the inner regions of haloes.995 Note that the GASn result in this panel agrees well with that in the left-hand panel. as does the additional GA3n profile plotted there for subhaloes with more than 30 yarticles (and so with miuc3010PAL an).," Note that the GA3n result in this panel agrees well with that in the left-hand panel, as does the additional GA3n profile plotted there for subhaloes with more than 30 particles (and so with $m_{\rm sub} > 3\times 10^{-6} M_{\rm halo}$ )."996 Although all these profiles are close to those. plotted in the upper yanels for mass-limitecl subhalo populations within haloes of much higher mass. they are nevertheless noticeably more concentrated.," Although all these profiles are close to those plotted in the upper panels for mass-limited subhalo populations within haloes of much higher mass, they are nevertheless noticeably more concentrated."997 This can be seen in Fig., This can be seen in Fig.998" LL. where we overplot he 30 particle limited subhalo number profile of €X3n and he mean profile for subhaloes with maa,c10LÀ, in our 15 clusters: the subhalo profiles are plotted: with svmbols."," 11, where we overplot the 30 particle limited subhalo number profile of GA3n and the mean profile for subhaloes with $m_{\rm sub}>10^{-4}M_{\rm halo}$ in our 15 clusters; the subhalo profiles are plotted with symbols."999 This suggests that as the density profile of the parent halo »ecomes more concentrated. so too does that of the subhalo population.," This suggests that as the density profile of the parent halo becomes more concentrated, so too does that of the subhalo population."1000 Note however. that the elect is much. weaker or the subhaloes than for the mass as a whole.," Note however, that the effect is much weaker for the subhaloes than for the mass as a whole."1001 Our subhalo number density. profiles are well fit by the following form: where. c is the distance to the host centre in units of roga. Gr) is the number of subhaloes within aw. UN. is the total number of subhaloes inside rogo. α=0.244. a=2. j—2.75. and e=Γιοι is the concentration of the host halo.," Our subhalo number density profiles are well fit by the following form: where, $x$ is the distance to the host centre in units of $r_{200}$ , $n(x$ ) is the number of subhaloes within $x$, $N$ is the total number of subhaloes inside $r_{200}$, $a=0.244$, $\alpha=2$, $\beta=2.75$, and $c=r_s/r_{200}$ is the concentration of the host halo."1002 The lines in Fig., The lines in Fig.1003 11 show the predications of this formula for €iX3n and forour 15 cluster haloes., 11 show the predications of this formula for GA3n and forour 15 cluster haloes.1004 Clearly.," Clearly,"1005conducting corona.,conducting corona.1006 It should be noted that application o other astrophysical svstenis is not trivial., It should be noted that application to other astrophysical systems is not trivial.1007 In SNR uodels. plasma expands iu a auch cooler. liuner aud ess conducting nediuu.," In SNR models, plasma expands in a much cooler, thinner and less conducting medium."1008 On the other haud. CAIE’s are uostlv observed iu the UV band. and are herefore to ve modeled as relatively cool perturbations. whereas the rou-confined frouts modeled here are at standard coronal eniperatures.," On the other hand, CME's are mostly observed in the UV band, and are therefore to be modeled as relatively cool perturbations, whereas the non-confined fronts modeled here are at standard coronal temperatures."1009 For the latter frouts to become proper CMTS. a mechauisii to thermally insulate the frouts frou he surrounding corona should be imvoked. so that they are ree to cool by radiation and to emit in the UV band (c Ciaravella et al.," For the latter fronts to become proper CME's, a mechanism to thermally insulate the fronts from the surrounding corona should be invoked, so that they are free to cool by radiation and to emit in the UV band (e.g. Ciaravella et al."1010 2001)., 2001).1011 Stellar fares are geucrally observed to evolve and decay onu time scales rangiug from several lows to davs. sugeesting large flaring reeious or even lack of magnetic confinement.," Stellar flares are generally observed to evolve and decay on time scales ranging from several hours to days, suggesting large flaring regions or even lack of magnetic confinement."1012 This work shows that the lvdrodvuamiuc evolution of flares occuring in non-confined. atmospheres leads invariably to a uch faster decay of the brightuess (on tie scales of very few nmunutes) after the heating phase. in a wide rauge of the plivsical paraincters.," This work shows that the hydrodynamic evolution of flares occurring in non-confined atmospheres leads invariably to a much faster decay of the brightness (on time scales of very few minutes) after the heating phase, in a wide range of the physical parameters."1013 One of the main implications is that the long duration of stellar flares is not indicativo of non-confned plasima., One of the main implications is that the long duration of stellar flares is not indicative of non-confined plasma.1014 Reversing the argument. it is highlv probable that the Observer long-lasting stellar N-ray flares involve forms of plasina confinement im closed coronal structures;," Reversing the argument, it is highly probable that the observed long-lasting stellar X-ray flares involve forms of plasma confinement in closed coronal structures."1015 We lave pointed out other characteristic features. with more or less diagnostical power. c.g. the svuchronous evolution of deusitv aud temperature. sunuumized im Table 2.. tha seen to indicate that N-rav observed stellar flares involve mostly plasma confined ii closed structures.," We have pointed out other characteristic features, with more or less diagnostical power, e.g. the synchronous evolution of density and temperature, summarized in Table \ref{tab:disc}, that seem to indicate that X-ray observed stellar flares involve mostly plasma confined in closed structures."1016 Although this work shuulates rather extreme conditions of fares occurriue in completely ron-confned. atinosplhieres. he result that long-lasting stellar flares are likely not occuring iu huge open strucures larecly notivates its validity.," Although this work simulates rather extreme conditions of flares occurring in completely non-confined atmospheres, the result that long-lasting stellar flares are likely not occurring in large open structures largely motivates its validity."1017 Frthermore this work predicts the possible existence of a new plienomenological class of events. characterized x peculiar light curves with a limited range of nne scales.," Furthermore this work predicts the possible existence of a new phenomenological class of events, characterized by peculiar light curves with a limited range of time scales."1018 Such sind of eveuts may be addressed by observations at Πο] sensitivity evels. which allow for lieh time resolutiou and high signal-to-noise ratio. such as those obained frou he current missions Chandra aud NATITNewton.," Such kind of events may be addressed by observations at high sensitivity levels, which allow for high time resolution and high signal-to-noise ratio, such as those obtained from the current missions Chandra and XMM-Newton."1019would be expected. to come from. relatively low. recdshilt galaxies (Cispertetal.2000).,would be expected to come from relatively low redshift galaxies \cite{gispert00}.1020.. The evolution in the high redshift regime is strongly. constrained by the CHltD/Sub-mm background at longer wavelengths where the ς slope is shallower than the average galaxy SED implying a strong contribution from the integrated light of high recshilt sources where the peak of the LR-SED is being sampled in the sub-mm waveband., The evolution in the high redshift regime is strongly constrained by the CIRB/Sub-mm background at longer wavelengths where the CIRB slope is shallower than the average galaxy SED implying a strong contribution from the integrated light of high redshift sources where the peak of the IR-SED is being sampled in the sub-mm waveband.1021 I should be noted that assuming an QO=] cosmology results in an extremely mariginal fit to the longer wavelength (sub-mm - mm) regime., It should be noted that assuming an $\Omega=1$ cosmology results in an extremely mariginal fit to the longer wavelength (sub-mm - mm) regime.1022 Constraints on the CLRB at shorter wavelengths come from DIRBE. LIES. (Matsumotoetal.2000).. galaxy source counts and TeV gamma rays (Stecker2000).. although the source count measurements and background. measurements are vet to converge on an agreed value probably due to the dillieultv in ascertaining the correction factor due to zodiacal light that peaks at —25tun (Ozernoy2000).," Constraints on the CIRB at shorter wavelengths come from DIRBE, IRTS \cite{mats00}, galaxy source counts and TeV gamma rays \cite{stecker00}, although the source count measurements and background measurements are yet to converge on an agreed value probably due to the difficulty in ascertaining the correction factor due to zodiacal light that peaks at $\sim$ $\umu$ m \cite{ozernoy00}."1023. Although the peak of the CIRB is well explained. by the strongly evolving ULEG population in this mioclel. the stronger contributor at sub-nun wavelengths are in fact the starburst galaxies.," Although the peak of the CIRB is well explained by the strongly evolving ULIG population in this model, the stronger contributor at sub-mm wavelengths are in fact the starburst galaxies."1024 Furthermore although much of the sub-mm background may be mace up from a population of ULIC: ealaxies (30-50% being resolved into individual galaxies with ο>J 2mgy). the surface density of the sub-mm population is significantly higher. (Hughes2000).," Furthermore although much of the sub-mm background may be made up from a population of ULIG galaxies $\%$ being resolved into individual galaxies with $S_{850}>2mJy$ ), the surface density of the sub-mm population is significantly higher \cite{hugh00}."1025.. Note that Eales et al., Note that Eales et al.1026 (2000) have pointed out that the SCUBA Iluxes may be upwardly biased by a factor of 1.4 due to source confusion and noise. resulting in a reduction of the resolved. CHAUD at S5Opum to 20/4," \shortcite{eales00} have pointed out that the SCUBA fluxes may be upwardly biased by a factor of 1.4 due to source confusion and noise, resulting in a reduction of the resolved CIRB at $\umu$ m to $\%$."1027 Such mocels of the CURB are also consistent with the Sb5ÜLun. source counts (see fig. 8))., Such models of the CIRB are also consistent with the $\umu$ m source counts (see fig. \ref{newcount}) ).1028" Phe ULICG component is completely dominant at bright ""κος down to 6m.Jv.", The ULIG component is completely dominant at bright fluxes down to $\sim$ 6mJy.1029 Below this limit the starburst galaxies begin to contribute more stronely. where at κου«2mJ/y they become the dominant source population.," Below this limit the starburst galaxies begin to contribute more strongly, where at $S_{850}<2mJy$ they become the dominant source population."1030 Peacock ct al., Peacock et al.1031 (2000) made a statistical analysis of the S50hun SCUBA map of the Llubble Deep Field (IDI) estimating that as much as 48% of the CIRB at S5OLUm could come from the sub-nim emission from the UV starforming galaxies with 30% from ULICG's etal.1998). and a further 10% from ACN CXlmainietal. 1999)., \shortcite{peacock00} made a statistical analysis of the $\umu$ m SCUBA map of the Hubble Deep Field (HDF) estimating that as much as $\%$ of the CIRB at $\umu$ m could come from the sub-mm emission from the UV starforming galaxies with $\%$ from ULIG's \cite{hugh98} and a further $\%$ from AGN \cite{almaini99}.1032. I the starburst population in the model is literally aken as the equivalent of the UV ealaxies in the LDIF then the model predictions slightly over predict this contribution from the starburst galaxies although the transition between starburst-LIG-ULICG is not clearly defined and therefore has a large scope for interpretation., If the starburst population in the model is literally taken as the equivalent of the UV galaxies in the HDF then the model predictions slightly over predict this contribution from the starburst galaxies although the transition between starburst-LIG-ULIG is not clearly defined and therefore has a large scope for interpretation.1033 Phe sub-muim source counts are also consistent with the results of Chapman et al., The sub-mm source counts are also consistent with the results of Chapman et al.1034 (2000) who carried oul a sub-mim survey of 33. Lyman break ealaxies (Aladauetal.1996). to —1.30iJy., \shortcite{chapman00} who carried out a sub-mm survey of 33 Lyman break galaxies \cite{madau96} to $\sim$ 1.3mJy.1035 The average lux of these sources was found to be Sis;0.6mJg implying that they would occupy the fainter end of the sub-mam source counts and would not significantly contribute to the brighter, The average flux of these sources was found to be $S_{850} \sim 0.6mJy$ implying that they would occupy the fainter end of the sub-mm source counts and would not significantly contribute to the brighter1036heir velocities.,their velocities.1037 The distribution of NoELGs (Fie. 5j) , The distribution of NoELGs (Fig. \ref{kernoem}) )1038is comparable to that derived from our auch larger photometric catalogue for all galaxies. indepeucently of spectral features and cluster moenbership (see Fie.," is comparable to that derived from our much larger photometric catalogue for all galaxies, independently of spectral features and cluster membership (see Fig."1039 l iu Slezak et al., 1 in Slezak et al.1040 1998): it is elongated alone PA~160° and shows a strong concentration around ΑΕ 85. a secondary peak towards the south cast coieciding with ABCG NT (sce for example Table 2 in Duet ct al.," 1998): it is elongated along $\sim 160^\circ$ and shows a strong concentration around ABCG 85, a secondary peak towards the south east coinciding with ABCG 87 (see for example Table 2 in Durret et al."1041 1998bj) aud au chhancement roughly at the position of ABCC 89 to the north west (as discussed by Duet et al., 1998b) and an enhancement roughly at the position of ABCG 89 to the north west (as discussed by Durret et al.1042 1998h)., 1998b).1043 The ealaxy distribution of ELGs (Fie. 6)), The galaxy distribution of ELGs (Fig. \ref{kerem}) )1044 is quite different: its peaks docs not coincide with ABCC 85. but is close to the position of ABCC 87: it shows a weak secondary maxima in the north-northeast direction aud clongatious along several PAs. all quite different from L160.," is quite different: its peak does not coincide with ABCG 85, but is close to the position of ABCG 87; it shows a weak secondary maximum in the north-northeast direction and elongations along several PAs, all quite different from $^\circ$."1045 The mean aud median velocities. as well as the velocity dispersions are quite different for ELCs and NoELGs: the iieun velocities are 15968 and 16627|... the median velocities ave 15701 aud 16732. aud the velocity dispersions are 1606 and 1109 or ELGs and NoELCs respectively. suggesting that the uorphologv-deusitv relation is coupled with kincmatic differences.," The mean and median velocities, as well as the velocity dispersions are quite different for ELGs and NoELGs; the mean velocities are 15968 and 16627, the median velocities are 15701 and 16732, and the velocity dispersions are 1606 and 1109 for ELGs and NoELGs respectively, suggesting that the morphology-density relation is coupled with kinematic differences."1046 Biviano et al. (, Biviano et al. (10471997) found differences in the average velocitics of ELGs and. NoELCs at a level larger han 20 oulv for 12 clusters out of their sample of 57: their interpretation was that in hese 12 clusters ELGs are a rou-virialized population falling onto the main cluster.,1997) found differences in the average velocities of ELGs and NoELGs at a level larger than $\sigma$ only for 12 clusters out of their sample of 57; their interpretation was that in these 12 clusters ELGs are a non-virialized population falling onto the main cluster.1048 In a much smaller sample of 6 clusters. Zabludotff Fraux (1993) also found a cifference in mean velocity between spirals aud carly type galaxies in 23 clusters: on the other haud. Mohr et al. (," In a much smaller sample of 6 clusters, Zabludoff Franx (1993) also found a difference in mean velocity between spirals and early type galaxies in 3 clusters; on the other hand, Mohr et al. ("10491996) found that in ABCC 576 ELCs and NoELGs had he same averageOo velocity. but with ELGs having a Luger velocity dispersion. as in ABCC 85.,"1996) found that in ABCG 576 ELGs and NoELGs had the same average velocity, but with ELGs having a larger velocity dispersion, as in ABCG 85."1050 Thevelocity distributions displaved in Figs., Thevelocity distributions displayed in Figs.1051 7 aud 8 were obtained simmlitancously using profile recoustructions based on a wavelet techuique and classical histograms., \ref{wavevnoem} and \ref{wavevem} were obtained simultaneously using profile reconstructions based on a wavelet technique and classical histograms.1052 We renuud the reader that the features obtained with the wavelet imiethod are significant at a 36 level above the noise (estimated at the smallest scale. see Fadda et al.," We remind the reader that the features obtained with the wavelet method are significant at a $\sigma$ level above the noise (estimated at the smallest scale, see Fadda et al."1053 1998)., 1998).1054 While the velocity distribution of NoELCs shows only oue peak around 16800 celose to the mean or median previously given. that of ELCs shows a peak at about 15300 aand a mich smaller one around 19250 1.," While the velocity distribution of NoELGs shows only one peak around 16800 close to the mean or median previously given, that of ELGs shows a peak at about 15300 and a much smaller one around 19250 ."1055 Therefore.," Therefore,"1056estimated richness close to the lower limit adopted or the inclusion of a candidate in the catalog (Αι30).,estimated richness close to the lower limit adopted for the inclusion of a candidate in the catalog $\Lambda_{cl}\sim30$ ).1057 A ligh-redshifts. only very rich chisters. probably with a large fraction of ellipticals. are detected in the two passbauds.," At high-redshifts, only very rich clusters, probably with a large fraction of ellipticals, are detected in the two passbands."1058 There are two such cases in the above table. but note that neither would have been imeluded as cluster candidates based on the V. detection alone.," There are two such cases in the above table, but note that neither would have been included as cluster candidates based on the $V$ detection alone."1059 Their appearance on the nuages strongly sugeests that both are at high recshifts., Their appearance on the images strongly suggests that both are at high redshifts.1060 However. since some galaxies are seen in the V oenages. either their matched filter redshitts ire overestimated or there are foreground concentrations cading to their detection iu the V. data.," However, since some galaxies are seen in the $V$ images, either their matched filter redshifts are overestimated or there are foreground concentrations leading to their detection in the $V$ data."1061 Ouly spectroscopic follow-up will be able to resolve such cases., Only spectroscopic follow-up will be able to resolve such cases.1062 The avadabilitv of data in two passbands cau. iu principle. provide an alternative wav of confirming cluster candidates and their estimated redshifts. based on detection of the sequeuce of cluster carly-type galaxies iu a C-À diagram.," The availability of data in two passbands can, in principle, provide an alternative way of confirming cluster candidates and their estimated redshifts, based on the detection of the sequence of cluster early-type galaxies in a C-M diagram."1063 Iu order to investigate this possibiliÁ [VI| diagramC» was produced for cach cluster candidate. showing all galaxies within a radius of 0.75 |! Mpe (Ij = 15 kun HL/Mpe) from the nominal cluster center.," In order to investigate this possibility, a $I$ $V-I$ ] diagram was produced for each cluster candidate, showing all galaxies within a radius of 0.75 $^{-1}$ Mpc $_0$ = 75 km $^{-1}$ /Mpc) from the nominal cluster center."1064 Fig., Fig.1065 2 shows four examples of such diagrsuus. for cluster candidates identified iu patches A and D. illustrating cases with estimated redshift in the range 0.2<+0.6.," \ref{fig:cm_diagram} shows four examples of such diagrams, for cluster candidates identified in patches A and B, illustrating cases with estimated redshift in the range $0.2 < z < 0.6$."1066 Also iudicated in the plot ave the values of in; aud the color of a tvplca elliptical (no-evolution) at the estimated redshift of the cluster. as derived from the matched filter.," Also indicated in the plot are the values of $m_I^*$ and the color of a typical elliptical (no-evolution) at the estimated redshift of the cluster, as derived from the matched filter."1067 At low redshift. the sequence of early-type galaxies is clearly visible. but at 220.5 the evidence for a C-M relation is. in nost cases. less compclline.," At low redshift, the sequence of early-type galaxies is clearly visible, but at $z \gsim 0.5$ the evidence for a C-M relation is, in most cases, less compelling."1068" Cousidering the combined patch A and B sample. oue finds that out of 35 clusters in the region of overlap of the V- aud Z-baud nuages. there are 19 with evidence for a C-M relation. with redshifts extending out to ><0.6,"," Considering the combined patch A and B sample, one finds that out of 35 clusters in the region of overlap of the $V$ - and $I$ -band images, there are 19 with evidence for a C-M relation, with redshifts extending out to $z \lsim 0.6$."1069 Furthermore. the redshift estimates based ou color and the matched filter secur to agree. dn most cases. within 0.1.," Furthermore, the redshift estimates based on color and the matched filter seem to agree, in most cases, within 0.1."1070 However. there are at least four cases where there is a strong sugecstion that the matched filter has overestimated the redshift.," However, there are at least four cases where there is a strong suggestion that the matched filter has overestimated the redshift."1071 Tn this letter 19 additional J cluster candidates detected in EIS patch D have been preseuted., In this letter 19 additional $I$ cluster candidates detected in EIS patch B have been presented.1072 Clusters have also been detected from the galaxy catalogs extracted frou the V images available in patches A and D. These clusters have been cross-identified with J detections. from which the following ecneral couclusious can be draw: The above results demonstrate the usefulness of V- observations for the robust detection of clusters.," Clusters have also been detected from the galaxy catalogs extracted from the $V$ images available in patches A and B. These clusters have been cross-identified with $I$ detections, from which the following general conclusions can be drawn: The above results demonstrate the usefulness of $V$ -band observations for the robust detection of clusters."1073 However. to extend the additional leverage provided bv the color information to redshifts huger than :~0.5 requires V exposures deeper than those obtained for LIS.," However, to extend the additional leverage provided by the color information to redshifts larger than $z \sim 0.5$ requires $V$ exposures deeper than those obtained for EIS."1074 This option should be considered for future surveys using wide-field cameras. iu particular in the Pilot Survey (Renzini 1998) to be conducted at the ESO 2.212 telescope.," This option should be considered for future surveys using wide-field cameras, in particular in the Pilot Survey (Renzini 1998) to be conducted at the ESO 2.2m telescope."1075The spectroscopic data are available for 14. CSS sources from the combined sample (Lable AL).,The spectroscopic data are available for 14 CSS sources from the combined sample (Table A1).1076 Phe ο/LEC: classification was possible for only 3 of them., The HEG/LEG classification was possible for only 3 of them.1077"08s00+472.. LEC"".", HEG*.1078 The image shows à compact source (~16kkpe) with a possible eastern. companion. 50 kkpc., The image shows a compact source $\sim$ kpc) with a possible eastern companion $\sim$ kpc.1079 The spectrum. shows a powerlaw continuum. ancl broad HIJA2800. emission. consistent with a QSO.," The spectrum shows a powerlaw continuum and broad $\lambda$ 2800 emission, consistent with a QSO."1080 The radio map shows compact source2001).. consistent with the QSO optical classification.," The radio map shows compact source, consistent with the QSO optical classification."1081"08s09+404.. LEC"".", HEG*.1082 The image shows à compact source (~12kkpe) with a possible tail towards W. consistent with the faint component in the radio map2001).," The image shows a compact source $\sim$ kpc) with a possible tail towards W, consistent with the faint component in the radio map."1083.1141--466.. The image shows an extended source (~20kkpe) wit 1a bright nucleus and bright emission knots ina very crowded Ποιά.1201+394., The image shows an extended source $\sim$ kpc) with a bright nucleus and bright emission knots in a very crowded field.1084. The image shows a compact source (~ kkpc). slightly elongated towards NW. consistent with the orientation of the radio structure2001).," The image shows a compact source $\sim$ kpc), slightly elongated towards NW, consistent with the orientation of the radio structure."10851241--411.. WEG., HEG.1086 The image shows an extended: source (~40 kkpc) with a bright nucleus., The image shows an extended source $\sim$ kpc) with a bright nucleus.10871343+386.. The optical data show a bright compact source (~20kkpe) wit 1a clear powerlaw and bright broad emission lines. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad emission lines, suggesting a QSO."10881445--410.. Phe image shows an extended. source with a bright nucleus., The image shows an extended source with a bright nucleus.1089 There is no clear evidence of the —20 kkpe radio structure in the optical image. 0801+303., There is no clear evidence of the $\sim$ kpc radio structure in the optical image. .1090. Lhe optical data show a bright compact source (730 kkpc) with a clear powerlaw and bright broad emission lines. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad emission lines, suggesting a QSO."10910853-4-291.. The optical data show a bright compact source (730 kkpc) with a clear powerlaw and bright broad emission lines. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad emission lines, suggesting a QSO."10921251+308.. The optical data show a bright compact source (30 κκρο with a clear powerlaw ancl bright broad WA2800 emission. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad $\lambda$ 2800 emission, suggesting a QSO."1093 The radio map shows a small source with a very structured. ~Skkpe long jet oriented SE-NW2006)., The radio map shows a small source with a very structured $\sim$ kpc long jet oriented SE-NW.1094.13154-396.. The optical data show a bright compact source (30 kkpe) wit1a clear powerlaw and bright broad emission ines. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad emission lines, suggesting a QSO."1095 The radio map shows a very small (~O.08 kkpe) jet oriented. EW2006)., The radio map shows a very small $\sim$ kpc) jet oriented EW.1096.1502+291.. The optical data show a bright compact source (30 kkpc) with a clear powerlaw and bright broad emission ines. suggesting a QSO.," The optical data show a bright compact source $\sim$ kpc) with a clear powerlaw and bright broad emission lines, suggesting a QSO."10971619+378.. The image shows a compact [faint source sugeesting a possible QSO., The image shows a compact faint source suggesting a possible QSO.1098 The spectrum shows only one clear line. with broad wings.," The spectrum shows only one clear line, with broad wings."1099 Lit is MgllH]A2800. then the source is at redshift z=1.2734.," If it is $\lambda$ 2800, then the source is at redshift $z=1.2734$."1100 However. this line could be alleeted by a possible break in the continuum.," However, this line could be affected by a possible break in the continuum."1101 It also shows no clear traces of the expected AGN powerlaw., It also shows no clear traces of the expected AGN powerlaw.1102" The radio map shows a 0.5"" (~4kkpe at 2=1.2734) jet oriented NIZ-SW. which does not show up in the optical image."," The radio map shows a $\arcsec$ $\sim$ kpc at z=1.2734) jet oriented NE-SW, which does not show up in the optical image."11031632+391.. The spectrum shows a clear powerlaw with bright broad. emission lines. consistent with a QSO.," The spectrum shows a clear powerlaw with bright broad emission lines, consistent with a QSO."1104 The image shows a two component source. separated. 230 kkpe.," The image shows a two component source, separated $\sim$ kpc."1105 The W component is compact ancl bright while. the [IZcomponent is fainter and. extended., The W component is compact and bright while the Ecomponent is fainter and extended.1106 The E component was classified. as à compact blue cluster by., The E component was classified as a compact blue cluster by.1107(1995).. he radio map shows a —Gkkpe jet oriented NE-SW at the. coordinates. of the compact source. which does not show up in the optical image.," The radio map shows a $\sim$ kpc jet oriented NE-SW at the coordinates of the compact source, which does not show up in the optical image."1108 Summing up. 9 sources (644)) are compact or unresolved. 5 sources (86%)) show extended emission. ancl 1 source (74) shows apossible merger or complex structure.," Summing up, 9 sources ) are compact or unresolved, 5 sources ) show extended emission, and 1 source ) shows apossible merger or complex structure."1109 Only 2 sources (14%)) show racio-optical alignment., Only 2 sources ) show radio-optical alignment.1110observations of several Landolt fields over many dozens of photometric nights.,observations of several Landolt fields over many dozens of photometric nights.1111 The estimated zero point errors are about 0.01 mae in all bands (Udalski et al., The estimated zero point errors are about 0.01 mag in all bands (Udalski et al.1112 1999b)., 1999b).1113 For more details about the instrumental svstem. observing. reduction and calibration procedures adopted in the course of the OGLE project the reader is relerred to the references cited above.," For more details about the instrumental system, observing, reduction and calibration procedures adopted in the course of the OGLE project the reader is referred to the references cited above."1114 The nem-infrared cata presented in (his paper were collected with the ESO NTT telescope on La Silla. equipped with the SOFT infrared camera.," The near-infrared data presented in this paper were collected with the ESO NTT telescope on La Silla, equipped with the SOFI infrared camera."1115" We used the Large Field setup with a field of view of 4.9 x 4.9 ατα αἱ a scale of 0,288. arcsec/pixel.", We used the Large Field setup with a field of view of 4.9 x 4.9 arcmin at a scale of 0.288 arcsec/pixel.1116" The gain and reaclout noise were 5.4 e/ADU and 0.4 ο, respectively."," The gain and readout noise were 5.4 e/ADU and 0.4 e, respectively."1117 The data were obtained in two observational programs. O74.D-0318(B). O074.D-0505(B) (PI: Pietrzvüsski) as part ol the Arancaria Project.," The data were obtained in two observational programs, 074.D-0318(B), 074.D-0505(B) (PI: Pietrzyńsski) as part of the Araucaria Project."1118 Single deep. J-band and A.-band observations of our target field were obtained under excellent seeing conditions during five different nights., Single deep $J$ -band and $K_s$ -band observations of our target field were obtained under excellent seeing conditions during five different nights.1119 On these nights. we also observed a large number (8-12) of photometric standard stars from (he UNIRT system (ILawarcen et al.," On these nights, we also observed a large number (8-12) of photometric standard stars from the UKIRT system (Hawarden et al."1120 2001) at a variety of airmasses and spanning a broad range in colors., 2001) at a variety of airmasses and spanning a broad range in colors.1121 To account for (he Irequent sky level variations in the infrared spectral region. especially in (he A.-band. (he observations were performed with a dithering technique.," To account for the frequent sky level variations in the infrared spectral region, especially in the $K_s$ -band, the observations were performed with a dithering technique."1122 For the A; and J-band observations we averaged over LO consecutive 10 second integrations (DITs) at any eiven pointing before moving the telescope to a randomly selected different position within a 25x25 arcsec square., For the $K_s$ and $J$ -band observations we averaged over 10 consecutive 10 second integrations (DITs) at any given pointing before moving the telescope to a randomly selected different position within a $\times$ 25 arcsec square.1123 Belween 15 and 25 such dithering positions were obtained. through ihe A.-band., Between 15 and 25 such dithering positions were obtained through the $K_s$ -band.1124 In the case of the J-band filter this number varied from 11 to 15., In the case of the $J$ -band filter this number varied from 11 to 15.1125 The reductions were performed in a similar manner to those described in Pietrzvisski and Gieren (2002)., The reductions were performed in a similar manner to those described in Pietrzyńsski and Gieren (2002).1126 The skv was subtracted [rom the images with a two-step process implying masking of the stars with the xdimsum ΗΑΕ package., The sky was subtracted from the images with a two-step process implying masking of the stars with the xdimsum IRAF package.1127 Then the individual images for each [ield ancl filter were flatfielded and stacked into a final composite image., Then the individual images for each field and filter were flatfielded and stacked into a final composite image.1128 The PSF photometry was carried out with the DAOPIIOT and ALLSTAR. programs., The PSF photometry was carried out with the DAOPHOT and ALLSTAR programs.1129 About 20-30 relatively bright and isolated stars were selected. visually ancl the first PSF model was derived [rom them., About 20-30 relatively bright and isolated stars were selected visually and the first PSF model was derived from them.1130 Following Pietrzvüsski. Gieren and Udalski (2002). we then ileratively improved the PSF model by subtracting all stars [rom their neighbourhood and re-caleulating the PSF model.," Following Pietrzyńsski, Gieren and Udalski (2002), we then iteratively improved the PSF model by subtracting all stars from their neighbourhood and re-calculating the PSF model."1131 After three such iterations no further improvement was noted. and the corresponding PSF model was adopted.," After three such iterations no further improvement was noted, and the corresponding PSF model was adopted."1132 In order to convert our PSF photometry (o (he aperture svstem. aperture corrections were derived for each frame.," In order to convert our PSF photometry to the aperture system, aperture corrections were derived for each frame."1133 This was done by using (he previously identilied candidates Lor PSF calculations after removing all nearby stars (hat could contaminate our photonmetry., This was done by using the previously identified candidates for PSF calculations after removing all nearby stars that could contaminate our photometry.1134 i>6. (277).,] $z> 6$ .1135. Iu particular. receutly argued that halos with circular velocities of up to ~30kins+ can be ploto-evaporated bv the UV background.," In particular, recently argued that halos with circular velocities of up to $\sim 30\ \rm km\,s^{-1}$ can be photo-evaporated by the UV background."1136 At the same time. the ionizing radiation may quickly dissociate molecular hydrogen. the ouly eficient coolaut for low-inctallicity gas in such halos. and prevent star formation before the eas is completely removed2).," At the same time, the ionizing radiation may quickly dissociate molecular hydrogen, the only efficient coolant for low-metallicity gas in such halos, and prevent star formation before the gas is completely removed."1137. The combined effect of these processes is likely to leave all DM. halos with masses <few10°ML. dark., The combined effect of these processes is likely to leave all DM halos with masses $\lesssim {\rm few}\times 10^9\ \rm M_{\odot}$ dark.1138 This Is consistent with current observational constraints which indicate that halos with AZ<Lot?AL. are virtually devoid of galaxies(7)., This is consistent with current observational constraints which indicate that halos with $M<10^{10}\ \rm M_{\odot}$ are virtually devoid of galaxies.1139. It is thus remarkable that the οσασα] mnasses of some of the Local Croup dwarfs are ouly ~10*NL. (2)...," It is thus remarkable that the dynamical masses of some of the Local Group dwarfs are only $\sim114010^7\ \rm M_{\odot}$ ."1141 How could such ealaxies form stars despite the suppressing processes listed above?, How could such galaxies form stars despite the suppressing processes listed above?1142 Oue possibility is that they manage to accrete a certain amount of gas before the Universe is reionized with the implicit assumption that tUs gas cali )o subsedquenutlv converted to stars., One possibility is that they manage to accrete a certain amount of gas before the Universe is reionized with the implicit assumption that this gas can be subsequently converted to stars.1143 Iowever. it is likely that gas cooling and star formation i such siunall svstenus is dneffücieut.," However, it is likely that gas cooling and star formation in such small systems is inefficient."1144 For example. cosinological simulations with self-cousisteut treatiueut of Πω clienistry and radiative trauster indicate that star formation is strougly suppressed in halos with masses ALX05s105AL. at all redshifts. even before relonization(7).," For example, cosmological simulations with self-consistent treatment of $_2$ chemistry and radiative transfer indicate that star formation is strongly suppressed in halos with masses $M\lesssim 5\times 10^8\ \rm M_{\odot}$ at all redshifts, even before reionization."1145. Iu addition. the galaxies may not be able to form sufficieutlv early to acercte the eas in the first place. if the power spectrum normalization is low or the Universe was reionized carly. as indicated by the first-vearWALAP results(??).," In addition, the galaxies may not be able to form sufficiently early to accrete the gas in the first place, if the power spectrum normalization is low or the Universe was reionized early, as indicated by the first-year results."1146. An alternativeproposal was recently sugeested bv aud corraborated bv ?.. who argued," An alternativeproposal was recently suggested by and corraborated by , who argued"1147expansion of the interaction region and magnetic reconnection has allowed the lield to penetrate through the entire region.,expansion of the interaction region and magnetic reconnection has allowed the field to penetrate through the entire region.1148 Thus the initial “wall” of tangle field wall is destroved and thermal conduction will be less inhibited than initially.," Thus the initial ""wall"" of tangle field wall is destroyed and thermal conduction will be less inhibited than initially."1149 We will check how accurately this proposed picture of destruction of field) wall is valid from analyzing our numerical simulations. and quantitatively discuss (he effects on the energy. transfer.," We will check how accurately this proposed picture of destruction of field wall is valid from analyzing our numerical simulations, and quantitatively discuss the effects on the energy transfer."1150 For our initial conditions. we set up an interface between hot and cold regions in mutual pressure equilibrium.," For our initial conditions, we set up an interface between hot and cold regions in mutual pressure equilibrium."1151 The temperature distribution on the horizontal (0) axis is given bv: in (he region 0c.r«I with Zy=100 in computational units., The temperature distribution on the horizontal $(x)$ axis is given by: in the region $0 < x < 1$ with $T_0 = 100$ in computational units.1152 This temperature prolile has a sharp gradient at c=0., This temperature profile has a sharp gradient at $x=0$.1153 The temperature distribution is plotted in Figure 2((a)., The temperature distribution is plotted in Figure \ref{fig02}( (a).1154 The region 0.4<r«0.5 is the interaction region we described in the previous section., The region $0.4 < x < 0.5$ is the interaction region we described in the previous section.1155 At the two side boundaries. (he temperature is set to be constant. and uniform across (he regions of each respective side of the box connecting to that side of the interaction region.," At the two side boundaries, the temperature is set to be constant, and uniform across the regions of each respective side of the box connecting to that side of the interaction region."1156 We are primarily interested in (he region of the box where (he heat transfer occurs aud noticeably evolves during (he simulation run time., We are primarily interested in the region of the box where the heat transfer occurs and noticeably evolves during the simulation run time.1157 This means we will mainly locus on the interaction region., This means we will mainly focus on the interaction region.1158 The horizontal length of the interaction region in the simulation domain is 0.1 in computational units., The horizontal length of the interaction region in the simulation domain is $0.1$ in computational units.1159 The thermal pressure is set to be in equilibrium over the entire box. Chat is," The thermal pressure is set to be in equilibrium over the entire box, that is"1160depends on a number of parameters. the most important of which are the ejectecl mass and its racial distribution.,"depends on a number of parameters, the most important of which are the ejected mass and its radial distribution."1161 The density profile of the gas in the newly born. SN is determined by the initial stellar structure. as modified by the explosion.," The density profile of the gas in the newly born SN is determined by the initial stellar structure, as modified by the explosion."1162 Numerical simulations of supernova explosions oduce density distributions that. during the [ree expansion hase. can be approximated by the functional formi psy=flee7 (seo ee. Chevalier Eransson 1994 and references herein).," Numerical simulations of supernova explosions produce density distributions that, during the free expansion phase, can be approximated by the functional form $\rho_{\rm SN}=f(v)t^{-3}$ (see e.g. Chevalier Fransson 1994 and references therein)."1163" The function. f(r) can in turn be represented. by a power-law in velocity. flr)xe""."," The function $f(v)$ can in turn be represented by a power-law in velocity, $f(v)\propto v^{-n}$."1164 To date. the best studied: case is that o£. SN. 1987X.. Modeling by Arnett (1988) and Shigevama Nomoto (1990) vield an almost lat inner powerlaw region. surrounded by a very steep outer »owerlaw profile. n9—10.," To date, the best studied case is that of SN 1987A. Modeling by Arnett (1988) and Shigeyama Nomoto (1990) yield an almost flat inner powerlaw region, surrounded by a very steep outer powerlaw profile, $n\sim 9-10$."1165" For normal abundances and at energies 10 keV. Chevalier Fransson (1994) estimate the optical depth at energy. ιο=££/(10keV) to he center of a supernova with a flat inner density. profile o be τςr£,DIS[FaMog where fsxsy ds the supernova enerey in units of LO erg. ancl Maio is the mass of the ejecta in units of LO AL..."," For normal abundances and at energies 10 keV, Chevalier Fransson (1994) estimate the optical depth at energy $E_{10}\equiv1166E/(10\;{\rm keV})$ to the center of a supernova with a flat inner density profile to be $\tau=\tau_s E_{10}^{-8/3}\; E_{\rm SN,116751}^{-3/2}\; M_{\rm ej, 10}^{5/2} t_{\rm yr}^{-2}$, where $E_{\rm SN,116851}$ is the supernova energy in units of $10^{51}$ erg, and $M_{\rm1169ej,10}$ is the mass of the ejecta in units of 10 $M_\odot$."1170" The constant 7, is ound to be 5.2 for a density profile with η=7 in the outer outs. and 4.7 for n=12."," The constant $\tau_s$ is found to be 5.2 for a density profile with $n=7$ in the outer parts, and 4.7 for $n=12$."1171 From these simple estimates. it can be seen that the SN would have to wait a decade or so before starting to become optically thin at the energies of interest.," From these simple estimates, it can be seen that the SN would have to wait a decade or so before starting to become optically thin at the energies of interest."1172 These estimates however do not account lor the fact that. if the SN harbors an energetic pulsar in its center. the pulsar itself will ionize a substantial fraction of the surrounding neutral material.," These estimates however do not account for the fact that, if the SN harbors an energetic pulsar in its center, the pulsar itself will ionize a substantial fraction of the surrounding neutral material."1173 Calculations of the ionization front of a pulsar in the interior of a voung SN were performed by Chevalier Fransson (1992)., Calculations of the ionization front of a pulsar in the interior of a young SN were performed by Chevalier Fransson (1992).1174" In the case of a flat. density profile in the inner region. ancl an outer density profile with powerkaw à?=9. they estimate that the ionization front reaches the edge of the constant density region. after NOEa timeη. /4,-.=10lofκ1ἐνboE1/35Maaoτουv1/2Dux where faut=Lt/l0tergs and f; is the fraction of the total rotational power that is converted in the form of ionizing radiation with a mean free path that is small compared to the supernova size."," In the case of a flat density profile in the inner region, and an outer density profile with powerlaw $n=9$, they estimate that the ionization front reaches the edge of the constant density region after a time $t_{\rm yr}=10\;t_01175f_i^{-1/3} \dot{E}_{\rm rot, 41}^{-1/3}\;M_{\rm ej, 10}^{7/6}\;E_{\rm1176SN, 51}^{-1/2}$, where $\dot{E}_{\rm rot, 41}\equiv\dot{E}_{\rm1177rot}/10^{41}\;{\rm erg}\;{\rm s^{-1}}$, and $f_i$ is the fraction of the total rotational power that is converted in the form of ionizing radiation with a mean free path that is small compared to the supernova size."1178 The constant fo depends of the composition of the core., The constant $t_0$ depends of the composition of the core.1179 For a hydrogen-dominated core. fo=1.64. for a helium-dominated core. fy=0.69. and for an oxvgen-dominated core fy=0.28.," For a hydrogen-dominated core, $t_0=1.64$, for a helium-dominated core, $t_0=0.69$, and for an oxygen-dominated core $t_0=0.28$."1180 Once the ionization [ront has reached the edge of the constant density region. the steep outer power-law. part of the density. profile is rapiclly ionized.," Once the ionization front has reached the edge of the constant density region, the steep outer power-law part of the density profile is rapidly ionized."1181" ""Therefore. depending on the composition aud total mass of the ejecta. an energetic pulsar can ionize the entire mass of the ejecta on a timescale between a few vears and a few tens of vears."," Therefore, depending on the composition and total mass of the ejecta, an energetic pulsar can ionize the entire mass of the ejecta on a timescale between a few years and a few tens of years."1182 This would clearly reduce the optical depth to the center of the remnant estimated above., This would clearly reduce the optical depth to the center of the remnant estimated above.1183 Given these considerations. in order to make predictions that are not as likely to be allected by opacity cllects. we also performed a Monte Carlo simulation of the compact remnant population for all the SNe with ages />10 vr. and another for allthe SNe with ages />30 vr.," Given these considerations, in order to make predictions that are not as likely to be affected by opacity effects, we also performed a Monte Carlo simulation of the compact remnant population for all the SNe with ages $t>10$ yr, and another for all the SNe with ages $t>30$ yr."1184 Since the opacity scales as £>. these subsets of objects are expected to be substantially less alfected by high optical depths to their inner regions.," Since the opacity scales as $t^{-2}$, these subsets of objects are expected to be substantially less affected by high optical depths to their inner regions."1185 The subsample of SNe with ages /10 vr contains +0 objects. while the subsample with ages /230 vr contains 2] SNe.," The subsample of SNe with ages $t>10$ yr contains 40 objects, while the subsample with ages $t>30$ yr contains 21 SNe."1186" The corresponding luminosity distributions (both measurements and limits) are shown in Figure 1 (middle and bottom panel respectively). together with the predictions of the adopted model (ACC initial period. cüstribution and Poe L,Esa correlation) for the luminosities of the pulsars associated with those SN samples."," The corresponding luminosity distributions (both measurements and limits) are shown in Figure 1 (middle and bottom panel respectively), together with the predictions of the adopted model (ACC initial period distribution and P02 $L_x-\dot{E}_{\rm rot}$ correlation) for the luminosities of the pulsars associated with those SN samples."1187 Given the uncertainties in the earlv-time optical depth. we consider the constraints derived from these subsamples (and especially the one with [o30 vr) more reliable.," Given the uncertainties in the early-time optical depth, we consider the constraints derived from these subsamples (and especially the one with $t>30$ yr) more reliable."1188 Furthermore. even independently of optical depth elfects that can bias the voungest members of the total sample. the subsamples of older SNe have on average lower luminosities. hence making the constraints on he model predictions more stringent.," Furthermore, even independently of optical depth effects that can bias the youngest members of the total sample, the subsamples of older SNe have on average lower luminosities, hence making the constraints on the model predictions more stringent."1189 In the following. when ecncralizing our study to derive limits on the allowed initial »eriod. distribution. we will use for our analysis only the subsets of older SNe.," In the following, when generalizing our study to derive limits on the allowed initial period distribution, we will use for our analysis only the subsets of older SNe."1190 In all three panels of Figure 1. the low bIuminosity tail of he simulation. accounting for 15% of the population. is dominated by the fraction of SNe whose compact remnants ave black holes. and for which we have assumed a luminosity ower than the lowest SN measurement/limit (~107 erg/s).," In all three panels of Figure 1, the low luminosity tail of the simulation, accounting for $\sim 15\%$ of the population, is dominated by the fraction of SNe whose compact remnants are black holes, and for which we have assumed a luminosity lower than the lowest SN measurement/limit $\sim 10^{35}$ erg/s)."1191 While it is possible that newly born DlIs could be accreting rom a fallback clisk and hence have luminosities as high as a few 107 erg/s our assumption of low Luminosity or them is the most conservative one for the analysis that we are performing. in that it allows us to derive the most stringent limits on the luminosity of the remaining remnant »opulation of neutron stars.," While it is possible that newly born BHs could be accreting from a fallback disk and hence have luminosities as high as a few $\times 10^{38}$ erg/s, our assumption of low luminosity for them is the most conservative one for the analysis that we are performing, in that it allows us to derive the most stringent limits on the luminosity of the remaining remnant population of neutron stars."1192 For these. the high-Iuminosity ail is ominated by the fastest. pulsars. those born with xriods of a few ms.," For these, the high-luminosity tail is dominated by the fastest pulsars, those born with periods of a few ms."1193 The magnetic fields. on the other vane. are in the bulk range of 10171045 G. Phe B field tail produces lower luminosities at. birth. while the ugh- [field tail will cause the pulsars to slow down on a imeseale smaller than the typical ages of the SNe in the sample.," The magnetic fields, on the other hand, are in the bulk range of $10^{12}-10^{13}$ G. The $B$ field tail produces lower luminosities at birth, while the $B$ field tail will cause the pulsars to slow down on a timescale smaller than the typical ages of the SNe in the sample."1194 Therefore. it is essentially the initial periods which xav a crucial role in determining the extent of the high-uminosity tail of the distribution.," Therefore, it is essentially the initial periods which play a crucial role in determining the extent of the high-luminosity tail of the distribution."1195" With the birth parameter distribution used here. we find that. out of the 10"". Monte Carlo realizations of the sample (for each of the three cases of Ες.) none of them preclicts pulsar luminosities compatible with the SN. N-rav These results point in the direction of initial periods of the pulsar population to be slower than the ms periods derived from some population svnthesis studies in the racio."," With the birth parameter distribution used here, we find that, out of the $10^6$ Monte Carlo realizations of the sample (for each of the three cases of Fig.1), none of them predicts pulsar luminosities compatible with the SN X-ray These results point in the direction of initial periods of the pulsar population to be slower than the ms periods derived from some population synthesis studies in the radio."1196 A number of other investigations in the last few vears. based on dilferent. methods of analysis of the radio sample with respect to ACC. have indeed come up to conclusions similar to ours.," A number of other investigations in the last few years, based on different methods of analysis of the radio sample with respect to ACC, have indeed come up to conclusions similar to ours."1197 The population svnthesis studies of Faucher-Giguere Ixaspi (2005) vielded a good fit to the data with the birth period described by a Gaussian with a mean period of 0.3 s and a spread of 0.15 sec., The population synthesis studies of Faucher-Giguere Kaspi (2005) yielded a good fit to the data with the birth period described by a Gaussian with a mean period of 0.3 s and a spread of 0.15 sec.1198 Similarly. the analvsis by Ferrario Wickramasinghe (2006) vielded a mean period. of 0.23 sec for a magnetic field of 107. G. We performed. Monte Carlo simulations of the X-ray pulsar population using the birth parameters derived in those studies above. and found," Similarly, the analysis by Ferrario Wickramasinghe (2006) yielded a mean period of 0.23 sec for a magnetic field of $10^{12}$ G. We performed Monte Carlo simulations of the X-ray pulsar population using the birth parameters derived in those studies above, and found"1199Measuring ooOealzwy properties as à function of their local cuviromment is a central task of observational extragalactic astronomy.,Measuring galaxy properties as a function of their local environment is a central task of observational extragalactic astronomy.1200 Doing so requires a measure of density., Doing so requires a measure of density.

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