ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 The resultscan be seen in ‘Table 3., The resultscan be seen in Table 3.3 Phe average error. (fromten hall skies) is below 12'4 down to a lux limit of 0.6 Jv (70 GLIz) and 0.4 Jy (100 CGLIz)., The average error (fromten half skies) is below $12\%$ down to a flux limit of $0.6$ Jy (70 GHz) and $0.4$ Jy (100 GHz).4 Phe determination of these cumulants would allow us, The determination of these cumulants would allow us5not only does the higher S/N data show smaller scatter (as expected. and indicated by the corresponcingly smaller error bars) but also exhibit Gaussian uncertainties. whereas the poorer quality data shows indications of a bias towards larger index values.,"not only does the higher S/N data show smaller scatter (as expected, and indicated by the correspondingly smaller error bars) but also exhibit Gaussian uncertainties, whereas the poorer quality data shows indications of a bias towards larger index values."6 This is particularly true of GC IX043 which appears as an outlier in al [our panels of Figure 5 and has the lowest S/N of the sample., This is particularly true of GC K043 which appears as an outlier in all four panels of Figure \ref{fig:ssp1} and has the lowest S/N of the sample.7 Concentrating solely on the ugher S/N data. we find that the models (plotted such that they are cllectively degenerate in age ancl exhibit ony a metallicity sequence) reproduce the cata very well.," Concentrating solely on the higher S/N data, we find that the models (plotted such that they are effectively degenerate in age and exhibit only a metallicity sequence) reproduce the data very well."8 “Phis is also the case Lor our composite metal-poor and metal-rich CC's., This is also the case for our composite metal-poor and metal-rich GCs.9 There is a sugeestion of a systematic olfset in the Mg6 Ales erid. since the composite GCs lic slightly above the &rids. although the olfset is not large.," There is a suggestion of a systematic offset in the Mg – $_2$ grid, since the composite GCs lie slightly above the grids, although the offset is not large."10" For comparison purposes. we also show the position of the central nuclear line-strengths of NGC 524 (through a standard 1.47. 4"" aperture). taken from Trager (1998)."," For comparison purposes, we also show the position of the central nuclear line-strengths of NGC 524 (through a standard $\arcsec \times $ $\arcsec$ aperture), taken from \citeANP{Trager98} (1998)."11 Two of the GCs in Figure 5 appear to have metallicities similar to that of NGC 524 itself., Two of the GCs in Figure \ref{fig:ssp1} appear to have metallicities similar to that of NGC 524 itself.12 In principle. α Κο ratios may be derived from these plots. but we defer further clisettssion Of this subject to the next section.," In principle, $\alpha$ /Fe] ratios may be derived from these plots, but we defer further discussion of this subject to the next section."13 For the rest of this analysis. we concenrate on the 14 GC's which have higher S/N. In Figure 6 we compare our data with the Alaraston&Thomas (2000) mocels. for the H (Pe5270 | be5335)/2. Mg». 111. Ley and L9i indices. and the combined. Mgle] index.," For the rest of this analysis, we concentrate on the 14 GCs which have higher S/N. In Figure \ref{fig:ssp2} we compare our data with the \citeANP{Maraston00} (2000) models, for the $\langle$ $\rangle$ (Fe5270 + Fe5335)/2, $_2$, $\beta$, $\gamma_A$ and $\delta_A$ indices, and the combined [MgFe] index."14 To assist in our interpretaticon of the erids. an additional metallicity interval at Fell] = O.S4 has been included in the mocdoels shown in Figure 6.. by linear interpolation between he Fel] = 1.35 and )33 lines.," To assist in our interpretation of the grids, an additional metallicity interval at [Fe/H] = –0.84 has been included in the models shown in Figure \ref{fig:ssp2}, by linear interpolation between the [Fe/H] = –1.35 and –0.33 lines."15 lt is clear that the individual data-points for the GC's populate a arge area of the SSP grid. parameter space., It is clear that the individual data-points for the GCs populate a large area of the SSP grid parameter space.16 In the lower tyvo panels of Figure ο the distribution in 112 is broadly οnsistent with the observational errors following the old-agec loci of the grids., In the lower two panels of Figure \ref{fig:ssp2} the distribution in $\beta$ is broadly consistent with the observational errors following the old-aged loci of the grids.17 The H4 index also behaves similarly: heavever itis clear that we can only measure L5 poorly., The $\delta_{\rm A}$ index also behaves similarly; however it is clear that we can only measure $\gamma_{\rm A}$ poorly.18 We have also plotted the nuclear line-streneths of NGC 5241πο taken from “Tragerctal. (1998). which indicate an old. metal-rich stellar svstem.," We have also plotted the nuclear line-strengths of NGC 524 itself, taken from \citeANP{Trager98} (1998), which indicate an old, metal-rich stellar system."19higher than our values. while the VCS 5 values for the power-law galaxies are uniformly. lower.,"higher than our values, while the VCS $\gamma'$ values for the power-law galaxies are uniformly lower."20 The dynamic range of the VCS +/ values is thus overall smaller than ours. and many of the VCS ! values place the galaxies in the intermediate zone between cores aud. power-laws.," The dynamic range of the VCS $\gamma'$ values is thus overall smaller than ours, and many of the VCS $\gamma'$ values place the galaxies in the intermediate zone between cores and power-laws."21 In a number ol cases. often where Ferrareseetal.(20060). claimed the existeuce of large uuclei. the agreement between the two samples is expecially poor. with the VCS 5 values falling in the range of core ealaxies [or svstems that we lad classified as power-laws.," In a number of cases, often where \citet{lf} claimed the existence of large nuclei, the agreement between the two samples is especially poor, with the VCS $\gamma'$ values falling in the range of core galaxies for systems that we had classified as power-laws."22 The ACS/WEC imagery used by the VCS study. aud the WEDPC2 and. WEPCI imagery that dominate the present sample provide essentially equivalent information., The ACS/WFC imagery used by the VCS study and the WFPC2 and WFPC1 imagery that dominate the present sample provide essentially equivalent information.23 A comparison of WFEPC2/PCI aud ACS/WEC PSFs shows that WEPC2 actually provides significantly better angular resolution (Figure 9))., A comparison of WFPC2/PC1 and ACS/WFC PSFs shows that WFPC2 actually provides significantly better angular resolution (Figure \ref{fig:psf}) ).24 The FWHM of the WEPC2/PCT F555W is 07061. while that of the ACS/WEC FL75W PSE is 07092. or larger.," The FWHM of the WFPC2/PC1 F555W is $0\asec061,$ while that of the ACS/WFC F475W PSF is $0\asec092,$ or larger."25 The use of (Fruchter&Hook2002) to rectify the ACS images slightly degrades its resolution further., The use of \citep{driz} to rectify the ACS images slightly degrades its resolution further.26 The present sample bas 19 galaxies at Virgo distance or closer. and a substantial number that are no more than umore clistaut: both the present aud VCS samples are probing the same plivsical scales in the galaxies.," The present sample has 49 galaxies at Virgo distance or closer, and a substantial number that are no more than more distant; both the present and VCS samples are probing the same physical scales in the galaxies."27 Reearcdless of the resolution dilfereuce between WEPC2 and ACS. PSE-deconvolution provides a modelindepeucent approach to correct ASTtnagery for the “wines” of the PSF outside theAST diffraction limit aud to isolate the profile analysis from the residual blurring interior to the diffraction limit that remains even after deconvolution.," Regardless of the resolution difference between WFPC2 and ACS, PSF-deconvolution provides a model-independent approach to correct imagery for the “wings” of the PSF outside the diffraction limit and to isolate the profile analysis from the residual blurring interior to the diffraction limit that remains even after deconvolution."28 The decouvolutiou tests presented in (1992b).. Laueretal. (1995).. Laueretal. (1993).. Restetal. (2001).. aud ave established that this methodology. works well forAST studies of brightness profiles: they also yield well-understood resolution limits.," The deconvolution tests presented in \citet{l92b}, \citet{l95}, \citet{l98}, \citet{rest}, and \citet{l05} have established that this methodology works well for studies of brightness profiles; they also yield well-understood resolution limits."29 The ACS images cau be decouvolved as well to correct for ferences between the ACS aud WEPC2 PSFs outside the resolution limit., The ACS images can be deconvolved as well to correct for differences between the ACS and WFPC2 PSFs outside the resolution limit.30 We demonustrate this with direct. Comparisons of decouvolved ACS and WEPC2 or WEPCI proliles for several galaxies --- COMMON., We demonstrate this with direct comparisons of deconvolved ACS and WFPC2 or WFPC1 profiles for several galaxies in common.31" Figure 10. compares proliles measured [roi PSF-cleconvolved ACS E175W. ""«drizzled"" images to those obtained from cdecouvolved WEPCI F555W images (Laueretal.1995). for eight galaxies iu common. and Crom decouvolved WEPC2 F552W images (Laueretal.2005) [or two galaxies in common."," Figure \ref{fig:acsdecon} compares profiles measured from PSF-deconvolved ACS F475W “drizzled"" images to those obtained from deconvolved WFPC1 F555W images \citep{l95} for eight galaxies in common, and from deconvolved WFPC2 F555W images \citep{l05} for two galaxies in common."32 We present more extensive comparisons between ACS aud WEPCI than between ACS aud WEPC2 because this is the more challenging test: the WEPCI imagery was blurred by a strongly aberrated PSF., We present more extensive comparisons between ACS and WFPC1 than between ACS and WFPC2 because this is the more challenging test; the WFPC1 imagery was blurred by a strongly aberrated PSF.33 Further. most of the power-law galaxies in coummon with the VCS sample were observed with WEDPCI rather than WEPC2: the steep central cusps in power-laws present a ereater challenge for deconvolution than do core galaxies.," Further, most of the power-law galaxies in common with the VCS sample were observed with WFPC1 rather than WFPC2; the steep central cusps in power-laws present a greater challenge for deconvolution than do core galaxies."34 The agreemente between all tliree cameras in all cases is excellent. with differences between the," The agreement between all three cameras in all cases is excellent, with differences between the"35"two clusters, its density would came out overestimated because background and foreground galaxies not physically associated with it would wrongly contribute to the In order to comply with criteria that are generally adopted in the literature (namely 1 Mpc, 1000 kms~!), so that the results from this work could be easily compared with other reference papers, but also keeping an eye on the physical structures that constitute the CS, we decided to evaluate the density adopting the following compromise strategy.","two clusters, its density would came out overestimated because background and foreground galaxies not physically associated with it would wrongly contribute to the In order to comply with criteria that are generally adopted in the literature (namely $1$ Mpc, $1000~$$\rm km s^{-1}$ ), so that the results from this work could be easily compared with other reference papers, but also keeping an eye on the physical structures that constitute the CS, we decided to evaluate the density adopting the following compromise strategy."36" First we compress the ""Fingers of God"" of Coma and Abell 1367 in the redshift space by assigning to galaxies belonging to the ""Fingers of God"" a velocity equal to the average velocity of the cluster, plus or minus a random Gaussian distributed AV comparable to the transverse size of the cluster on the plane of the sky, assuming that clusters have approximately a spherical shape in 3-D. (We assume that the transverse size is 2 deg and 1 deg for Coma and A1367, corresponding to 248 and 124 kms""! respectively)."," First we compress the ""Fingers of God"" of Coma and Abell 1367 in the redshift space by assigning to galaxies belonging to the ""Fingers of God"" a velocity equal to the average velocity of the cluster, plus or minus a random Gaussian distributed $\Delta V$ comparable to the transverse size of the cluster on the plane of the sky, assuming that clusters have approximately a spherical shape in 3-D. (We assume that the transverse size is 2 deg and 1 deg for Coma and A1367, corresponding to 248 and 124 $\rm km s^{-1}$ respectively)."37keV aud Ady=2.82 keV aud a metal abuudauce Z—0.21 provides a good fit to the integrated spectrum.,"keV and $kT_2=382.82$ keV and a metal abundance $Z = 0.21$ provides a good fit to the integrated spectrum."39(Pitjeva2009) and INPOP (Fiengaοἱal.20098) ephemericdes by including some vears of continuous radiometric ranging data to Cassini in addition to data of several types spanning the last century. (Pitjeva2008:Fiengaοἱal.2009b) both are non-zero al a statistically significant level (80 aud 1.26. respectively) and thev are compatible each other since their difference is equal to 4+10 mas |.,"\citep{Pit09} and INPOP \citep{INPOP}40 ephemerides by including some years of continuous radiometric ranging data to Cassini in addition to data of several types spanning the last century, \citep{Pit08,Fie09}41 both are non-zero at a statistically significant level $3\sigma$ and $1.2\sigma$, respectively) and they are compatible each other since their difference is equal to $4\pm 10$ mas $^{-1}$."42 At themonent*.. no corrections A estimated with the DE ephemerides by NASA JPL are available.," At the, no corrections $\Delta\dot\varpi$ estimated with the DE ephemerides by NASA JPL are available."43 lorio(2009a) unsuccessfully examined several possible ανασα] explanations in terms of both mundane. standard: Newlonian/relativislic gravitational phivsies ancl οἱ modified models of gravity.," \citet{Ior09} unsuccessfully examined several possible dynamical explanations in terms of both mundane, standard Newtonian/relativistic gravitational physics and of modified models of gravity."44 Anyway. further analvses of extended data sets from Cassini with different dynamical force models are required to firmly. establish the existence of the anomalous perihelion precession of Saturn as a genuine physical effect.," Anyway, further analyses of extended data sets from Cassini with different dynamical force models are required to firmly establish the existence of the anomalous perihelion precession of Saturn as a genuine physical effect."45 llere we will show that the existence of a localized distant body. (planet X/Nemesis). modeled in neither EPM nor INPOP ephenmerides. is a good candidate to explain a secular peribelion precession of Saturn having the characteristics of(2): indeed. contrary (o a massive ring usually adopted to model the action of the minor asteroids and of the Trans Neptunian Objects (TNOs). it vields a retrograde secular perihelion precessions and (he constraints on its distance for different postulated values of ils mass are consistent with several theoretical predictions put forth to accommodate some features of the Edgeworth-Ixuiper bell (Κακα&Mukai2008).," Here we will show that the existence of a localized distant body (planet X/Nemesis), modeled in neither EPM nor INPOP ephemerides, is a good candidate to explain a secular perihelion precession of Saturn having the characteristics of: indeed, contrary to a massive ring usually adopted to model the action of the minor asteroids and of the Trans Neptunian Objects (TNOs), it yields a retrograde secular perihelion precessions and the constraints on its distance for different postulated values of its mass are consistent with several theoretical predictions put forth to accommodate some features of the Edgeworth-Kuiper belt \citep{Lyk}."46. Concerning Nemesis. it would be an undiscovered stellar companion of the Sun which. moving along a highlv elliptical!.. would periodically disturb the Oort cloud being responsible of the periodicity of about 26 Alvr in extinction rates on the Earth over the last 250 Myr (Whitmire&al.1984): : (he Nemesis hypothesis has also been used to explain (he measurements ol the ages of 155 lunar spherules from the Apollo 14 site (Muller2002).," Concerning Nemesis, it would be an undiscovered stellar companion of the Sun which, moving along a highly elliptical, would periodically disturb the Oort cloud being responsible of the periodicity of about 26 Myr in extinction rates on the Earth over the last 250 Myr \citep{Nem1,Nem2}; the Nemesis hypothesis has also been used to explain the measurements of the ages of 155 lunar spherules from the Apollo 14 site \citep{Mul02}."47. See e.g. for huther details., See e.g. \citet{IorNem} for further details.48 Interestingly. such a proposed explanation of the anomalous perihelion precession of Saturn in terms of pointlike dark matter is. to a certain extent. to be considered as degenerate since also the MOdified Newtonian Dynamics (MOND) (Milgrom1953) predicts certain subtle effects in the planetary region of the solar svstem," Interestingly, such a proposed explanation of the anomalous perihelion precession of Saturn in terms of pointlike dark matter is, to a certain extent, to be considered as degenerate since also the MOdified Newtonian Dynamics (MOND) \citep{Mil83} predicts certain subtle effects in the planetary region of the solar system"49fibre radius of 1.25 aresec. and under the assumption of an exponential surface brightness profile. a mean scale length of 2.9 aresec. a mean effective radius of 3.6 arcsec (Jerjen Dressler 19972) and for the case that the fibre had been ideally centred onto the galaxy.,"fibre radius of 1.25 arcsec, and under the assumption of an exponential surface brightness profile, a mean scale length of 2.9 arcsec, a mean effective radius of 3.6 arcsec (Jerjen Dressler 1997a) and for the case that the fibre had been ideally centred onto the galaxy."50 Two fields were excluded from the analysis because of their much lower yield due to problems with the autoguider., Two fields were excluded from the analysis because of their much lower yield due to problems with the autoguider.51 The limit in surface brightness which could be reached here confirms that projects aiming at the systematic measurement of redshifts for dwarf galaxies in nearby clusters are feasible. even with a multi-fibre instrument.," The limit in surface brightness which could be reached here confirms that projects aiming at the systematic measurement of redshifts for dwarf galaxies in nearby clusters are feasible, even with a multi-fibre instrument."52 The only requirement is a large field of view., The only requirement is a large field of view.53 0.5cm Correction for the earth motion was then carried out. leading to heliocentric corrected redshifts for both the cross-correlation and the emission-line redshifts.," 0.5cm Correction for the earth motion was then carried out, leading to heliocentric corrected redshifts for both the cross-correlation and the emission-line redshifts."54 For 32 galaxies the redshifts obtained here could be compared to those obtained previously by one of us (Stein 1996)., For 32 galaxies the redshifts obtained here could be compared to those obtained previously by one of us (Stein 1996).55 Note that cross-correlation errors for the brightest galaxies in Stein (1996) might be slight underestimations. due to the fact that the template had been constructed using a sample of these same bright galaxies.," Note that cross-correlation errors for the brightest galaxies in Stein (1996) might be slight underestimations, due to the fact that the template had been constructed using a sample of these same bright galaxies."56 Thus. the scaling factor between internal (Tonry Davis 1979) and external errors is determined using only galaxies in the calibration sample with errors of 20 or above.," Thus, the scaling factor between internal (Tonry Davis 1979) and external errors is determined using only galaxies in the calibration sample with errors of 20 or above."57 With an error scaling factor of 1.7 the differences between redshifts in both datasets are perfectly consistent with the resulting external errors., With an error scaling factor of 1.7 the differences between redshifts in both datasets are perfectly consistent with the resulting external errors.58 À zero-point shift of 26 was then applied to the cross-correlation data., A zero-point shift of 26 was then applied to the cross-correlation data.59 Only for six galaxies it was possible to obtain both cross-correlation and emission-line redshifts., Only for six galaxies it was possible to obtain both cross-correlation and emission-line redshifts.60 Again. the distribution of differences is in good agreement with the expectations. given the uncertainties in both measurements.," Again, the distribution of differences is in good agreement with the expectations, given the uncertainties in both measurements."61 Finally. a weighted mean of emission and absorption (cross-correlation) redshifts was taken.," Finally, a weighted mean of emission and absorption (cross-correlation) redshifts was taken."62 We measured redshifts for 115 galaxies., We measured redshifts for 115 galaxies.63 The data are listed in 55a. 5b. and 5e combined with data from the CCC such as morphological type. total apparent D mmagnitude. ane SB.g.," The data are listed in 5a, 5b, and 5c combined with data from the CCC such as morphological type, total apparent $B$ magnitude, and $_{\rm eff}$."64 Galaxies with no previous redshift measurement are subdivided into two lists according to cluster members 55a) and background objects 55b)., Galaxies with no previous redshift measurement are subdivided into two lists according to cluster members 5a) and background objects 5b).65 55e contains the data for galaxies which already had a redshift measured and are confirmed cluster members., 5c contains the data for galaxies which already had a redshift measured and are confirmed cluster members.66 Among our galaxies there are 101 with velocities smaller than 5414ο., Among our galaxies there are 101 with velocities smaller than 5414.67.. This velocity corresponds to the 3 upper limit for the velocity distribution of Cen45 (LCD) and was used as a cut-off to discriminate between cluster members and background galaxies., This velocity corresponds to the $\sigma$ upper limit for the velocity distribution of Cen45 (LCD) and was used as a cut-off to discriminate between cluster members and background galaxies.68 Actually. the remaining 14 velocities lie between [00000 and 0000 which makes a separation unambiguous.," Actually, the remaining 14 velocities lie between 000 and 000 which makes a separation unambiguous."69 Our cluster sample was complemented by 19 redshifts taken from the literature (DCL: LC: Stein 1996)., Our cluster sample was complemented by 19 redshifts taken from the literature (DCL; LC; Stein 1996).70 Finally. the dataset upon which the following analysis is based consists of redshifts for 120 cluster members located in the central cluster area.," Finally, the dataset upon which the following analysis is based consists of redshifts for 120 cluster members located in the central cluster area."71 A weighted average was taken in case of nultiple redshifts. after homogenization of zero-point-shifts and sealing of errors among the four sources.," A weighted average was taken in case of multiple redshifts, after homogenization of zero-point-shifts and scaling of errors among the four sources."72 The completeness of this data set is very high with respect to apparent magnitude and effective surface brightness., The completeness of this data set is very high with respect to apparent magnitude and effective surface brightness.73 Redshifts are available for 96 percent of all known cluster galaxies brighter than 7217.5., Redshifts are available for 96 percent of all known cluster galaxies brighter than $B_{\rm T}$ =17.5.74 At the limit of Br=18.5 our redshift sample is still complete to 78 percent.," At the limit of $B_{\rm75T}$ =18.5 our redshift sample is still complete to 78 percent."76" This latter magnitude corresponds to Mp=— 15.3. assuming a cluster distance modulus of (0.M(44,233.79 (Jerjen Dressler 1997b) and including an extinction. term of 4520.42."," This latter magnitude corresponds to $M_{B_{\rm T}}=-$ 15.3, assuming a cluster distance modulus of $(m-M)_{\rm Cen}$ =33.79 (Jerjen Dressler 1997b) and including an extinction term of $A_{\rm B}$ =0.42."77" For $B,.g the completeness levels are 93 percent at Baaresee?. and 78 percent at Baaresce 7. respectively."," For $_{\rm eff}$ the completeness levels are 93 percent at $\,B$ $^{-2}$ and 78 percent at $B$ $^{-2}$, respectively."78 55a and 5b contain data for 50 galaxies with new measured velocities among which 36 galaxies are cluster members according to our selection criterion., 5a and 5b contain data for 50 galaxies with new measured velocities among which 36 galaxies are cluster members according to our selection criterion.79 Their Hubble type mixture is (1/3/11/21)., Their Hubble type mixture is (1/3/11/21).80 The composition of the background sample is (0/3/6/5)., The composition of the background sample is (0/3/6/5).81 For the latter sample the morphological information may not necessarily be true any more because the classification had been done under the assumption of cluster membership., For the latter sample the morphological information may not necessarily be true any more because the classification had been done under the assumption of cluster membership.82 The new measured redshifts shall be used to estimate the accuracy of the morphological based cluster membership for dwarf galaxies in the CCC., The new measured redshifts shall be used to estimate the accuracy of the morphological based cluster membership for dwarf galaxies in the CCC.83 For this purpose we compare in 22 the number of cluster members and background objects for the two dwarf families individually., For this purpose we compare in 2 the number of cluster members and background objects for the two dwarf families individually.84 In 11. we give the three membership classes as listed in the CCC., In 1 we give the three membership classes as listed in the CCC.85 22 and 3 give the observed numbers of late-type dwarfs of a particular membership class divided into members (m) and background (b) according to their redshifts., 2 and 3 give the observed numbers of late-type dwarfs of a particular membership class divided into members (m) and background (b) according to their redshifts.86 From 22 and 3 we derive the fraction of real cluster members in 44 which can be compared to the percentages listed in 11., From 2 and 3 we derive the fraction of real cluster members in 4 which can be compared to the percentages listed in 1.87 The same analysis is done for the early-type dwarfs in the columns 5-7., The same analysis is done for the early-type dwarfs in the columns 5–7.88 As can be seen. the numbers in 11. 4 and 7 are in good agreement.," As can be seen, the numbers in 1, 4 and 7 are in good agreement."89 The case is slightly, The case is slightly90is that. higher observation density makes detection of high eccentricity planets more likely.,is that higher observation density makes detection of high eccentricity planets more likely.91 The subsets of 1179949 show significant. variations in Di(e;). but is there a difference between stars?," The subsets of 179949 show significant variations in $D^\prime_{\mathrm{int}}(e_i)$, but is there a difference between stars?"92" Figure 14 shows Di,(e;) for each of the three objects. 1179949. 220782 and 338382."," Figure \ref{fig:eselfunc1} shows $D^\prime_{\mathrm{int}}(e_i)$ for each of the three objects, 179949, 20782 and 38382."93 The shape of the curves. while in general decreasing at higher eccentricities. is dilleren [or each object.," The shape of the curves, while in general decreasing at higher eccentricities, is different for each object."94 For example. when e;0.1 1179049 has the lowest fraction of planets redetected. however when e;2OS 1 iw the highest.," For example, when $e_i\le 0.1$ 179949 has the lowest fraction of planets redetected, however when $e_i\ge 0.8$ it has the highest."95" Thus data sampling ai quality are fundamental to the selection elfects. present in planet search observations and. a simple parametrisation of the detectability of cxoplanct parameters using ""whole-of-survey” metrics e.g. 2 done.", Thus data sampling and quality are fundamental to the selection effects present in planet search observations and a simple parametrisation of the detectability of exoplanet parameters using ``whole-of-survey'' metrics – e.g. \citet{Cumming04} –.96 As an example. consider the case of 1179949.," As an example, consider the case of 179949."97" Of. the three sets. of observations we discuss here. the 1179949 data have the highest median measurement uncertainty ty, and one might naively expect it’s detectabilities to. he the lowest."," Of the three sets of observations we discuss here, the 179949 data have the highest median measurement uncertainty $^{-1}$ ), and one might naively expect it's detectabilities to be the lowest."98 However. the observation density (equal to the observation time-span/number of observations or 7NTZN) is the highest at ddays/epoch. whieh should counteract the first elect to some degree.," However, the observation density (equal to the observation time-span/number of observations or $\Delta T/N$ ) is the highest at days/epoch, which should counteract the first effect to some degree."99 It is not intuitively clear how to parametrise ancl compare the cetectability of the 1179949 observations. with. for example. that of the 338382. observations which have lower observation density but also lower median measurement. uncertainty without the simulations we have carried. out in this study.," It is not intuitively clear how to parametrise and compare the detectability of the 179949 observations, with, for example, that of the 38382 observations – which have lower observation density but also lower median measurement uncertainty – without the simulations we have carried out in this study."100 Therefore carrying out simulations on a star-by-star basis is the way to understand the selection ellects in Doppler velocity. planet searches., Therefore carrying out simulations on a star-by-star basis is the way to understand the selection effects in Doppler velocity planet searches.101 In the previous section. we examined DX(0;).. the detectability at each ο. which is at all 2? and AJ;.," In the previous section, we examined $D^\prime_{\mathrm{int}}(e_i)$ , the detectability at each $e_i$, which is at all $P_i$ and $M_i$."102 This is fine [or the case of these simulations. where we know the input parameter valuespriori.," This is fine for the case of these simulations, where we know the input parameter values."103" Llowever. as this is never the case for actual Doppler planet data it is useful to consider our detectability at each 6,,: Le. the eccentricities rather than the input eccentricities."," However, as this is never the case for actual Doppler planet data it is useful to consider our detectability at each $e_m$; i.e. the eccentricities rather than the input eccentricities."104 We call this quantity Dii0)., We call this quantity $D^\prime_{\mathrm{int}}(e_m)$.105" Leis determined by counting the number of correct detections i.c. false positives are excluded: in. equally spaced bins of e,,. normalised by the number of —-simulations in cach of bin."," It is determined by counting the number of correct detections – i.e. false positives are excluded – in equally spaced bins of $e_m$ , normalised by the number of simulations in each of bin."106We] now compare thetwo quantities.S (ο).' and,"We now compare thetwo quantities, $D^\prime_{\mathrm{int}}(e_m)$ and"107as a polytrope and that pοςT?[sin(za)/(12)]?/ (see Eq. [6]]).,as a polytrope and that $\rho \propto T^{3} \propto [sin(\pi x)/(\pi x)]^{3/4}$ (see Eq. \ref{eq:rho}] ]).108" However, the external layers of a SN typically have a steeper power-law distribution caused by the"," However, the external layers of a SN typically have a steeper power-law distribution caused by the"109finding extended the underprediction of high-mass subhalos to an underprediction of luminous satellites. implying that the MW-MCs system is unusual.,"finding extended the underprediction of high-mass subhalos to an underprediction of luminous satellites, implying that the MW-MCs system is unusual."110 Similarly. Okamotoetal.(2010) explored a range of feedback models to add galaxies to some of the high resolution Aquarius halos (Springeletal.2008).," Similarly, \citet{Okamoto10} explored a range of feedback models to add galaxies to some of the high resolution Aquarius halos \citep{Springel08}."111. It was again difficult to readily reproduce halos with lummosities as bright as the MCs., It was again difficult to readily reproduce halos with luminosities as bright as the MCs.112 Having looked in detail at a handful of simulated objects. however. this work indicates that there may be significant halo-to-halo scatter in the number of such massive objects.," Having looked in detail at a handful of simulated objects, however, this work indicates that there may be significant halo-to-halo scatter in the number of such massive objects."113 The idea of intrinsic scatter in the subhalo population was expanded on in Ishiyamaetal. (2009a)., The idea of intrinsic scatter in the subhalo population was expanded on in \cite{Ishiyama09}.114". While this work did not concentrate specifically on MC-like subhalos. they did consider the range in number of subhalos with vj,sub!Verhos20.1."," While this work did not concentrate specifically on MC-like subhalos, they did consider the range in number of subhalos with $v_{max,sub}/v_{max,host} > 0.1$."115 This work showed an extremely large variation (20-60) in the number of such massive subhalos a galaxy-sized halos would host., This work showed an extremely large variation (20-60) in the number of such massive subhalos a galaxy-sized halos would host.116 In contrast to the semi-analytic modeling just discussed. Libeskindetal.(2007) used hydrodynamic simulations to model the luminosity functions for the satellites around MW-like host halos.," In contrast to the semi-analytic modeling just discussed, \cite{Libeskind07} used hydrodynamic simulations to model the luminosity functions for the satellites around MW-like host halos."117 Their simulation identified 9 MW-like central galaxies and found that they. on average. have 1.6 satellites brighter than My.2—16 and a third of them had satellites with luminosities comparable to the LMC.," Their simulation identified 9 MW-like central galaxies and found that they, on average, have 1.6 satellites brighter than $M_V =118-16$ and a third of them had satellites with luminosities comparable to the LMC."119 The recent Millennium-II and Bolshor simulations (Boylan-Kolchinetal.2010:Klypin2010) have. for the first time. allowed us to probe cosmological volumes to understand the predictions for the satellite populations of MW-like halos.," The recent Millennium-II and Bolshoi simulations \citep{BoylanKolchin10,Klypin10} have, for the first time, allowed us to probe cosmological volumes to understand the predictions for the satellite populations of MW-like halos."120 The properties of these simulations are summarized in Table 1., The properties of these simulations are summarized in Table 1.121 Boylan-Kolchinetal.(2010) (hereafter BKΙΟ) quantified the likelihood for 10'7M halos to host massive satellite galaxies in the Millennium-II ..simulation. finding that subhalos similar to the MCs are quite rare.," \citet{BoylanKolchin10}122 (hereafter BK10) quantified the likelihood for $10^{12} \msol$ halos to host massive satellite galaxies in the Millennium-II simulation, finding that subhalos similar to the MCs are quite rare."123 In this paper. we expand on this work by making similar measurements. for the Bolshor simulation. which used WMAP7 cosmological parameters. and using an abundance matching technique to make detailed comparisons between the Bolshoi predictions and the measurements from Liuetal.(2010).," In this paper, we expand on this work by making similar measurements for the Bolshoi simulation, which used WMAP7 cosmological parameters, and using an abundance matching technique to make detailed comparisons between the Bolshoi predictions and the measurements from \cite{Liu10}."124. Our goal is to understand just how well reproduces the statistical properties of bright satellites., Our goal is to understand just how well reproduces the statistical properties of bright satellites.125 Note that this is the reverse question from the one that was asked in a companion paper. Bushaetal.(2010)..," Note that this is the reverse question from the one that was asked in a companion paper, \cite{Busha10c}."126 That work assumed a satellite population and asked what the implications were for the properties of the host halo. including its mass.," That work assumed a satellite population and asked what the implications were for the properties of the host halo, including its mass."127 Here. we assume a host halo mass and ask about the implications for the subhalo population.," Here, we assume a host halo mass and ask about the implications for the subhalo population."128 There ts no reason for both questions to give the same answer: while Bushaetal.(2010) showed that a halo which hosts two MC-like satellites most likely has a mass near 1.2«10'M.. there is no reason to assume that a typical 1.2«I0'7M.. halo will have the MC:an as satellites.," There is no reason for both questions to give the same answer: while \cite{Busha10c} showed that a halo which hosts two MC-like satellites most likely has a mass near $1.2 \times 10^{12} \msol$, there is no reason to assume that a typical $1.2 \times 10^{12}\msol$ halo will have the MCs as satellites."129 We begin by giving an overview of the Bolshoi simulation ins 2.. and then investigate the properties of massive dark matter satellites around dark matter host halos in $3.. focusing on the mass ranges for hosts and satellites that are most relevant to the MW system.," We begin by giving an overview of the Bolshoi simulation in \ref{sec:sims}, and then investigate the properties of massive dark matter satellites around dark matter host halos in \ref{sec:halos}, focusing on the mass ranges for hosts and satellites that are most relevant to the MW system."130 The analysis here is similar to that of BKIO., The analysis here is similar to that of BK10.131 In $4.. we assign galaxy luminosities to our suite of dark matter halos and extend the results for a sample with similar selection cuts as for observations.," In \ref{sec:luminosities}, we assign galaxy luminosities to our suite of dark matter halos and extend the results for a sample with similar selection cuts as for observations."132 In this way. we are able to make detailed comparisons to the observational work of (Liuetal.2010.hereafterLIO) concerning the satellite population around MW-magnitude galaxies.," In this way, we are able to make detailed comparisons to the observational work of \cite[][hereafter L10]{Liu10} concerning the satellite population around MW-magnitude galaxies."133 Section 4.4 gives the results of this analysis — see especially Figure 8.., Section \ref{sec:comparisons} gives the results of this analysis — see especially Figure \ref{fig:nsats_obs_sim}.134 Finally. in 35.. we expand this study to include the satellite population of a more general distribution of hosts. and $6 summarizes our conclusions.," Finally, in \ref{sec:generalproperties}, , we expand this study to include the satellite population of a more general distribution of hosts, and \ref{sec:conclusions} summarizes our conclusions."135 Throughout this paper. we adopt the convention /=0.7 (the value that was used in the Bolshoi simulation) when reporting values from either simulations or observations.," Throughout this paper, we adopt the convention $h =1360.7$ (the value that was used in the Bolshoi simulation) when reporting values from either simulations or observations."137 We use the dark matter halos identified in the Bolshoi simulation (Klypinetal.2010:Trujillo-Gomez2010).," We use the dark matter halos identified in the Bolshoi simulation \citep{Klypin10, TrujilloGomez10}."138" This simulation modeled a 250 /r!Mpe comoving box using cosmological parameters similar to those derived by WMAP7 (Komatsuetal.2010): ©,,=0.27. O4=0.73. oy=0.82. i20.95. and h20.7."," This simulation modeled a 250 $\hmpc$ comoving box using cosmological parameters similar to those derived by WMAP7 \citep{Komatsu10}: $\Omega_m = 0.27$ , $\Omega_{\Lambda} = 0.73$, $\sigma_8 = 0.82$, $n=0.95$, and $h = 0.7$."139 The simulation volume contains 2048? particles. each with a mass of 1.35«10°7!M.. and was run using the ART code (Kravtsovetal.1997).," The simulation volume contains $2048^3$ particles, each with a mass of $1.35 \times 10^8~\hinv\msol$ and was run using the ART code \citep{Kravtsov97}."140. 180 snapshots from the simulation were saved and analyzed., 180 snapshots from the simulation were saved and analyzed.141 One of the unique aspects of this simulation ts the high level of spatial resolution employed. allowing objects to be resolved down to a physical scale of 1 A7'kpe.," One of the unique aspects of this simulation is the high level of spatial resolution employed, allowing objects to be resolved down to a physical scale of 1 $\hkpc$."142 This gives us excellent ability to track halos as they merge with and are disrupted by larger objects. allowing us to track them even as they pass near the core of the host halo.," This gives us excellent ability to track halos as they merge with and are disrupted by larger objects, allowing us to track them even as they pass near the core of the host halo."143 A summary of these simulation parameters is presented in Table Τ.., A summary of these simulation parameters is presented in Table \ref{table:simulations}.144 Because we discuss our work in the context of the BKΙΟ results. we also present the same parameters for the Millennium II simulation (Boylan-Kolchinetal. 2009).. on which the ΒΚ10 results were based.," Because we discuss our work in the context of the BK10 results, we also present the same parameters for the Millennium II simulation \citep{BoylanKolchin09}, on which the BK10 results were based."145 Halos and subhalos were identified using the BDM algorithm (Klypin&Holtzman1997)., Halos and subhalos were identified using the BDM algorithm \citep{Klypin97}.146.. The algorithm identifies maxima in. the density field and examines the neighboring region to identify bound particles., The algorithm identifies maxima in the density field and examines the neighboring region to identify bound particles.147 In this way. it treats both halos and subhalos identically.," In this way, it treats both halos and subhalos identically."148 Subhalos are just identified as objects living within the virial radius of a larger objects., Subhalos are just identified as objects living within the virial radius of a larger objects.149" Because of the high level of mass and spatial resolution. BDM results in a halo catalog that is complete down to a maximum circular velocity v4,250 km s!, where Vinay=Max(GmnDE This corresponds to à virial mass of roughly 10/57!M..."," Because of the high level of mass and spatial resolution, BDM results in a halo catalog that is complete down to a maximum circular velocity $\vmax = 50$ km $^{-1}$, where $\vmax = {\rm max}\left(\sqrt{{GM(<r)}/{r}}\right). $ This corresponds to a virial mass of roughly $10^{10}\hinv\msol$."150 When BDM halos are identified. the ids of their 50 most bound particles are also stored to assist in producing merger trees.," When BDM halos are identified, the ids of their 50 most bound particles are also stored to assist in producing merger trees."151 As discussed in Section. 4.2.. in order to assign galaxy luminosities to dark matter halos. we need to track the histories of dark matter substructures.," As discussed in Section \ref{sec:sham}, in order to assign galaxy luminosities to dark matter halos, we need to track the histories of dark matter substructures."152 This is done using merger trees created from the 180 snapshots of the Bolshoi simulation., This is done using merger trees created from the 180 snapshots of the Bolshoi simulation.153 The detailed algorithm for creating the merger trees is deseribed in Behroozietal.(2010a)., The detailed algorithm for creating the merger trees is described in \cite{Behroozi10b}.154. Briefly. the algorithm works by first linking halos across time steps by tracking the 50 most bound particles of each halo.," Briefly, the algorithm works by first linking halos across time steps by tracking the 50 most bound particles of each halo."155 Some halos will not have any of their 50 most bound particles identified at a later timestep (e.g.. very massive objects in which the 50 most bound particles change rapidly through stochastic processes). while some will have their particles distributed to multiple halos.," Some halos will not have any of their 50 most bound particles identified at a later timestep (e.g., very massive objects in which the 50 most bound particles change rapidly through stochastic processes), while some will have their particles distributed to multiple halos."156 The algorithm corrects for this by running a simple N-body calculation on the locations and masses of all halos in the simulation to predict where each halo should wind up at the next time step., The algorithm corrects for this by running a simple N-body calculation on the locations and masses of all halos in the simulation to predict where each halo should wind up at the next time step.157 Using this information. it is possible to link halos across multipletime steps with very high accuracy.," Using this information, it is possible to link halos across multipletime steps with very high accuracy."158 We begin by investigating the properties of dark matter satellites around dark matter hosts in the Bolshoi simulation. focusing on the mass range for hosts and satellites that is relevant to the MW system.," We begin by investigating the properties of dark matter satellites around dark matter hosts in the Bolshoi simulation, focusing on the mass range for hosts and satellites that is relevant to the MW system."159 In $3.1.. we consider trends with," In \ref{sec:massdep}, , we consider trends with"160"velocity of NM d. rus, = 175.8 ((CGathier et al.","velocity of M $-$ 4, $v_{\rm LSR}$ = $-$ 175.8 (Gathier et al."161 1983) it is unlikely to be a foreground object., 1983) it is unlikely to be a foreground object.162 In 2 we give the input values for the observables used. for the modelling. together with the resulting moclel oedicetions.," In \ref{ratag:inp} we give the input values for the observables used for the modelling, together with the resulting model predictions."163 As can be seen [rom this table. not all the ines present in the spectra are predicted byCLOUDY.. most notably the higher Balmer lines of hydrogen anc several velit lines.," As can be seen from this table, not all the lines present in the spectra are predicted by, most notably the higher Balmer lines of hydrogen and several helium lines."164 Also the clement chlorine is not. included. in he code., Also the element chlorine is not included in the code.165 The resulting physical parameters for the nebulae are given in 3.., The resulting physical parameters for the nebulae are given in \ref{ratag:phys}.166 The hydrogen density shown in this able is the constant density within the ssphere., The hydrogen density shown in this table is the constant density within the sphere.167 The PNe in our sample all have nearly the same medium excitation class., The PNe in our sample all have nearly the same medium excitation class.168 “This probably is partially a result of our selection criterion that the nebulae should have been detected by iin the 12 bband (criterion 1)., This probably is partially a result of our selection criterion that the nebulae should have been detected by in the 12 band (criterion 1).169 Old bulge PNe. having a high excitation class and cool dust. might have insullicient 12 flux to be detected by45.," Old bulge PNe, having a high excitation class and cool dust, might have insufficient 12 flux to be detected by."170. In the rest of this section each of the PNe in our sample will be discussed individually. with special emphasis on the problems encountered during the modelling.," In the rest of this section each of the PNe in our sample will be discussed individually, with special emphasis on the problems encountered during the modelling."171 ‘Two lines were omitted from the list of observables because of the following reasons., Two lines were omitted from the list of observables because of the following reasons.172 First theLei A4686 line was omitted. because the Dux ratio given by RPDAL is quite high. indicative of a high stellar temperature.," First the 4686 line was omitted, because the flux ratio given by RPDM is quite high, indicative of a high stellar temperature."173 However. the rest of he observational data are not consistent. with such a high stellar temperature.," However, the rest of the observational data are not consistent with such a high stellar temperature."174 Also. this line is listed in 3 of UPDA. but is not present in their 1.," Also, this line is listed in 3 of RPDM, but is not present in their 1."175 Webster (1988. WSS) took a spectrum of this PN. and she didn't report he detection of this line.," Webster (1988, W88) took a spectrum of this PN, and she didn't report the detection of this line."176 She should however have detected a line of the strength mentioned by PDA., She should however have detected a line of the strength mentioned by RPDM.177 Tylenda et al., Tylenda et al.178 (1994) list an upper limit of5 for the intensity of this line.," \cite{ty94}179 list an upper limit of 5 for the intensity of this line."180 In view of these uncertainties we eecided to omit this line., In view of these uncertainties we decided to omit this line.181 Since RPDAL included this line in their modelling. this probably explains the higher stellar temperature they. obtain.," Since RPDM included this line in their modelling, this probably explains the higher stellar temperature they obtain."182 The fitting of then1] A4363 line was also problematic., The fitting of the 4363 line was also problematic.183 The observed. flux was far too low to be consistent. with the electron temperature predicted by our model., The observed flux was far too low to be consistent with the electron temperature predicted by our model.184 Since the electron temperature derived. from theui] line ratio is much higher (and more consistent with the value determined by our model). and also because the11] A4363 line is much stronger in the spectrum of WSs (however not as strong as predicted by our model). we decided that its value was too uncertain and omitted it from the input.," Since the electron temperature derived from the line ratio is much higher (and more consistent with the value determined by our model), and also because the 4363 line is much stronger in the spectrum of W88 (however not as strong as predicted by our model), we decided that its value was too uncertain and omitted it from the input."185 The intensity of the line seems quite high. and is not fitted well.," The intensity of the line seems quite high, and is not fitted well."186 The discrepancy is too large to be attributed. to measurement errors. rence this might indicate that the spectrum has not been sullicientIy. derecdened.," The discrepancy is too large to be attributed to measurement errors, hence this might indicate that the spectrum has not been sufficiently dereddened."187 Phere is however no evidence from he fits to the other lines to support this suspicion., There is however no evidence from the fits to the other lines to support this suspicion.188 Our model gives a very small inner radius. also resulting in à very high ionization parameter.," Our model gives a very small inner radius, also resulting in a very high ionization parameter."189 This is caused by the ugh 112 oover 25 lux ratio. which might indicate the presence of hot dust.," This is caused by the high 12 over 25 flux ratio, which might indicate the presence of hot dust."190 see also the discussion in Paper I. The spectrum is fitted: quite well. but there is. slight discrepancy for theOu] A4959 and. A5007 lines.," See also the discussion in Paper I. The spectrum is fitted quite well, but there is slight discrepancy for the 4959 and 5007 lines."191 This is caused by theΟΠ 3727 doublet. which is not fitted well.," This is caused by the 3727 doublet, which is not fitted well."192 Phe latter doublet usually has a larger uncertainty due to extinction and detector insensitivity., The latter doublet usually has a larger uncertainty due to extinction and detector insensitivity.193 This PN has the highest 47 of all PNe in our sample., This PN has the highest $\chi^2$ of all PNe in our sample.194" ""This is mainly caused by the weak lines. which might indicate that his spectrum has a lower signal-to-noise when compared to he other spectra."," This is mainly caused by the weak lines, which might indicate that this spectrum has a lower signal-to-noise when compared to the other spectra."195 RPDAL do not list error margins for their ine [lux ratios. so we had to assume reasonable values.," RPDM do not list error margins for their line flux ratios, so we had to assume reasonable values."196 The model is not able to fit the 225 flux. which is very high compared both to the 12 aand GO flux.," The model is not able to fit the 25 flux, which is very high compared both to the 12 and 60 flux."197 A possible explanation could be the presence of a 30 deust. feature in the spectrum. (Lloare 1990)., A possible explanation could be the presence of a 30 dust feature in the spectrum (Hoare 1990).198 This would imply that the nebula is carbon-rich. since this feature has only been observed in carbon-rich nebulae.," This would imply that the nebula is carbon-rich, since this feature has only been observed in carbon-rich nebulae."199 The central star has spectral tvpe OF (Aller Ixeyes LOST. AINST).," The central star has spectral type Of (Aller Keyes 1987, AK87)."200 The large cilference between the optical diameter. of (CXcker et al., The large difference between the optical diameter of (Acker et al.201 1992) and the radio diameter of, 1992) and the radio diameter of202but functionally equivalent expression forjj have been used in the literature 1996).,"but functionally equivalent expression for$\mu$ have been used in the literature \citep[Milgrom 1983b, 1984;][]{sanders96}."203. Currently the most used functiou. also adopted here. is: which works well in describing galactic rotation curves aud has the advantage of being simple.," Currently the most used function, also adopted here, is: which works well in describing galactic rotation curves and has the advantage of being simple."204 Solving for g the true acceleration of gravity reads This expression allows the calculation of the true gravitational acceleration provided. the Newtonian one is known., Solving for $g$ the true acceleration of gravity reads This expression allows the calculation of the true gravitational acceleration provided the Newtonian one is known.205" Before proceediug to compute ος. it is worth noticiug that if 7 is the true value of M/L of an object. aud 7, is the one derived using Newton's law to convert. accelerations to lasses. then g/g,=7/7."," Before proceeding to compute $\gn$, it is worth noticing that if $\tau$ is the true value of M/L of an object, and $\taun$ is the one derived using Newton's law to convert accelerations to masses, then $g/\gn \equiv \taun / \tau$."206 ludeed. imagine an acceleration g is meastwed at distance r [οιan object of ltuninosity L.," Indeed, imagine an acceleration $g$ is measured at distance $r$ froman object of luminosity $L$ ."207" In MOND we interpret this g as due to a mass M=gr?/G[]7L. where the term [] is the [actor g/g, delined in eq.3."," In MOND we interpret this $g$ as due to a mass $M=g r^2/G[]=\tau L$, where the term $[]$ is the factor $g/\gn$ defined in eq.3."208" Ou the other hand. according to Newton's law we infer a nass M,=qi?/Gn L."," On the other hand, according to Newton's law we infer a mass $\Mn=g r^2/G = \taun L$ ."209" Then the ratio M,/M=7,/r{|g/gy.", Then the ratio $\Mn/M=\taun/\tau = [] = g/\gn$.210" Thus. any dilfereuce between g and g, will appear as a variation of τ.To investigate these effects ou elliptical galaxies. E shall take as represeutative of each galaxy the acceleration of gravity gay at ο. computed setting 7 from the stellar population typical of the ealaxies aud replacing the gravitational radius with ο."," Thus, any difference between $g$ and $\gn$ will appear as a variation of $\tau$ .To investigate these effects on elliptical galaxies, I shall take as representative of each galaxy the acceleration of gravity $g_N$ at $r_e$, computed setting $\tau$ from the stellar population typical of the galaxies and replacing the gravitational radius with $r_e$."211 Iu abseuce of dark matter both assumptions are valid., In absence of dark matter both assumptions are valid.212 The mass within rc is set to be ;c=O.11 of the total mass. as appropriate for galaxies following the ;A£1 gle Vaucouleurs law (Young1976).," The mass within $r_e$ is set to be $x=0.41$ of the total mass, as appropriate for galaxies following the $r^{1/4}$ de Vaucouleurs law \citep{young76}."213. Under these assumptions. the acceleration ol gravity at re becomes: where the numerical coustaut is correct for 7 aud. L in solar units. aud r« in pc.," Under these assumptions, the acceleration of gravity at $r_e$ becomes: where the numerical constant is correct for $\tau$ and $L$ in solar units, and $r_e$ in pc."214 Combinine eq., Combining eq.215" 3 aud 1 oue can compute g/g, for all galaxies.", 3 and 4 one can compute $g/\gn$ for all galaxies.216 Belore proceediug with the calculation of gy. E have to set the value of 7.," Before proceeding with the calculation of $g_N$ , I have to set the value of $\tau$ ."217 SinceE am assuming Newtons law [ails belowaj. 1uasses derived usiug auy dyuuulcal mocel based on this law cau obviously not be used here.," SinceI am assuming Newton's law fails below$a_0$, masses derived using any dynamical model based on this law can obviously not be used here."218The spectra of the 2 erating settiugs covering the IT band are shown iu Fig. l..,"The spectra of the 2 grating settings covering the H band are shown in Fig. \ref{hspec},"219 aud of the 3 erating settiugs covering the I band in Fig. 2.., and of the 3 grating settings covering the K band in Fig. \ref{kspec}.220 The K baud top panel shows the stroug > and IL euission lines known from previous low-resolution spectroscopy (Castro-Tirado 11996)., The K band top panel shows the strong $\gamma$ and I emission lines known from previous low-resolution spectroscopy (Castro-Tirado 1996).221 Similarly. we see Brackett series in cluission frou 11-1 4) to 15-1 A) iu the IT baud.," Similarly, we see Brackett series in emission from 11-4 $\eta$ ) to 15-4 $\lambda$ ) in the H band."222 We do uot fine eudssiou as reported |w Castro-Tirado ((1996). Eikeuberry (1998) anc ((2000). supporting their conclusion that this is a variable feature probadv related to the N-rav state aud jet ejection activity.," We do not find emission as reported by Castro-Tirado (1996), Eikenberry (1998) and (2000), supporting their conclusion that this is a variable feature probably related to the X-ray state and jet ejection activity."223 We note that during our 1999 ISAAC observation of GRS 1915105 was in a state of low activity at N-ravs and radio. though the time from the last radio flare and towards the rext radio flare were different to the ((2000) observation.," We note that during our 1999 ISAAC observation of GRS 1915+105 was in a state of low activity at X-rays and radio, though the time from the last radio flare and towards the next radio flare were different to the (2000) observation."224 Both. the II as we Las lires are clearly rexἼνος. vaving FWIIM — 1015 aat a resolution of 5 (α ithe ID bou.," Both, the H as well as lines are clearly resolved, having FWHM $\sim$ 10–15 at a resolution of 5 (in the H band)."225 Iftus were die to rotational Doppler adenine. it would correspond to a velocity of e sm 200300 kiu/s. In nost cases. these lines are οΗΝ but have a ceutra depression.," If this were due to rotational Doppler broadening, it would correspond to a velocity of $v$ $i$ $\sim$ 200–300 km/s. In most cases, these lines are not gaussian, but have a central depression."226 Caven the Ct that he inchlation of the binary svsteni ds 7TUdeg#2¢ce (MiraIC Rodneuez 1991) oue indeed may alicipate a double-li edsrape., Given the fact that the inclination of the binary system is $i \sim 70\deg\pm2\deg$ (Mirabel Rodriguez 1994) one indeed may anticipate a double-lined shape.227 We do not fd P Cre xofiles iu he Dr5 aud II as reported by ((2000)., We do not find P Cyg profiles in the $\gamma$ and I as reported by (2000).228 Iu acdcition. we fiuc fortjo first time several absortion ines which allow us to make a rough ideutification of the donor in the CRS 19151105 binary.," In addition, we find for the first time several absorption lines which allow us to make a rough identification of the donor in the GRS 1915+105 binary."229 Iu tιο WN uid we clearly ideutifv 300 absorption baud heads characteristic of a low te3uperature (TTOU| Is) star (c.g. Wletmmanun Hall 1986)., In the K band we clearly identify $^{12}$ CO absorption band heads characteristic of a low temperature $T< 7000$ K) star (e.g. Kleinmann Hall 1986).230 Though weak. we also identity the CO (2.0 yand CO (3.1) trausilolis. indicating a bhunuimositv class IIT or brighter (e.g. Wallace Iüukle 1997).," Though weak, we also identify the $^{13}$ CO (2,0) and $^{13}$ CO (3,1) transitions, indicating a luminosity class III or brighter (e.g. Wallace Hinkle 1997)."231 We also ileutifv t1 Na. doublet (2.2062L/2.20897 gnu). aud possibly the Ca triplet you). ALT (9.0088 μπι). aud the MgII doublet au) m absorption.," We also identify the Na doublet (2.20624/2.20897 $\mu$ m), and possibly the Ca triplet $\mu$ m), I (2.10988 $\mu$ m) and the I doublet $\mu$ m) in absorption."232 Note that the CN doublet jiu). which In superelauts is more proniuneut than Al/Mg. is not detected.," Note that the CN doublet $\mu$ m), which in supergiants is more prominent than Al/Mg, is not detected."233 Iu the II baud. we iceutiv Mell jn) hough 2CO (LL) may also conrbute) “CO (6.3) and CO (8.5) iu a ratio which is consistent with MEI standards (Mover 11998). axl II (16718.9/16750.6 jn).," In the H band, we identify I $\mu$ m) (though $^{12}$ CO (4,1) may also contribute), $^{12}$ CO (6,3) and $^{12}$ CO (8,5) in a ratio which is consistent with MK standards (Meyer 1998), and I (16718.9/16750.6 $\mu$ m)."234 Comparing the 2.32.1 in ea spectrum from 20/21 July 1999 with hat taken on 21/25 July 2000 (after heliocentric correction). we find tvat the CO band head systems are shifted hy 60 lau's yolativ to cach other.," Comparing the 2.3–2.4 $mu$ m spectrum from 20/21 July 1999 with that taken on 24/25 July 2000 (after heliocentric correction), we find that the CO band head systems are shifted by 60 km/s relativ to each other."235 The easiest interpretalon is Dopder moion. aud therefore midicates that the CO absorp10 ris indeed of photospheric origin aud not due to absorptio Lina static. cold. eieuustellar ecu.," The easiest interpretation is Doppler motion, and therefore indicates that the CO absorption is indeed of photospheric origin and not due to absorption in a static, cold, circumstellar medium."236 Tus. we couchde tji we have identified the «onor in CRS 1915105. iux that it is a late-tvpe. K-M giant.," Thus, we conclude that we have identified the donor in GRS 1915+105, and that it is a late-type, K-M giant."237 We have tried to coufirm the DIuunositv class more quantitatively by using the veiline-independent indicator r=loglLEW{302.O0)ΓΗ(Na)EW(Ca))| (Raiiirrez 11997).," We have tried to confirm the luminosity class more quantitatively by using the veiling-independent indicator $~~~~~r = \log [EW(^{12}{\rm CO} (2,0))/(EW({\rm Na}) + EW({\rm Ca}))]$ rez 1997)."238 Because oftie low significance of tie Ca triplet. our measurement has a laree error: r=O0.25+0.20.," Because of the low significance of the Ca triplet, our measurement has a large error: $r = 0.25 \pm 0.20$."239 This value falls in between the ranges covered by dwirfs 0.2z5e 50.0) and elants 15r SO.6) (Ramirrez 119097)., This value falls in between the ranges covered by dwarfs $r$ 0.0) and giants $r$ 0.6) rez 1997).240 T1ο ratio of equivalent wiIths of 12CO to PCO which depends ou Iuuinositv class (Campbell 119903. has heen measured for the seven transitions covered. (ower panel of Fig. 2))," The ratio of equivalent widths of $^{12}$ CO to $^{13}$ CO which depends on luminosity class (Campbell 1990), has been measured for the seven transitions covered (lower panel of Fig. \ref{kspec}) )"241 to 3+l. again supporting a giaut classification.," to $\sim 3 \pm 1$, again supporting a giant classification."242 Usiug the IT band spectra. we also considered the veilme-iuxependent temperature/Iuuinositv discriminant EW(OT 1.6901 jun)/EW(Mg 1.5765. pon} vs. EW(CO 1.6610 µια. | CO L6187 μι)ΤΝλος 1.5765 flu) as proposed w Mover (1998).," Using the H band spectra, we also considered the veiling-independent temperature/luminosity discriminant EW(OH 1.6904 $\mu$ m)/EW(Mg 1.5765 $\mu$ m) vs. EW(CO 1.6610 $\mu$ m + CO 1.6187 $\mu$ m)/EW(Mg 1.5765 $\mu$ m) as proposed by Meyer (1998)."243" The OII line is. unfortunately, only marginally detected. aud therefore only the huninosity class caunot be : ↸⊳∪∐↴∖↴⊓⋅⋜"," The OH line is, unfortunately, only marginally detected, and therefore only the luminosity class cannot be constrained."244↧∐∐∖≺↧∙↽∕∏∐∖↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖↸∖↴∖↴↑∐⊔⋜↧↑↸∖⋅↖⇁↕↸∖↕≼⇂↴∖↴∿↓≺∖∩↸⇅⋅ ⋅ ∖⊐⋃⋃EN Is. which would sugseest a late-C 00r dX spectral type (IIoudashelt," The temperature estimate yields $\sim$ $^{+200}_{-500}$ K, which would suggest a late-G or K spectral type (Houdashelt"245trend Chat we see in all the models of starless cores. ie. the first scenario. (an example is shown in Fig. 9)),"trend that we see in all the models of starless cores, i.e. the first scenario, (an example is shown in Fig. \ref{mod_pre}) )"246 is that. once the final density is reached (at (5.27x 109 ves). the molecular abundances of the observed species evolve since thev stabilize. changing by no more (han a few hundredthis with respect to the final value: the higher the depletion efficiency (he shorter the time needed to reach the stable state.," is that, once the final density is reached (at $\times$ $^6$ yrs), the molecular abundances of the observed species evolve since they stabilize, changing by no more than a few hundredths with respect to the final value: the higher the depletion efficiency the shorter the time needed to reach the stable state."247 Another general trend is that the column density ol has a double peak. the first maximum is reached soon after the final density is reached. the second is the final value.," Another general trend is that the column density of has a double peak, the first maximum is reached soon after the final density is reached, the second is the final value."248 This double peak behavior was also found by other authors (e.g. Gwenlanοἱal. 2000))., This double peak behavior was also found by other authors (e.g. \citealt{gwenlan00}) ).249 In Fig., In Fig.250 9. we show the evolution of the column densities of the observed species in a starless core (ie. in the first scenario) with a density of 5x10! and à FR=0.2. that corresponds toa percentage of CO on grains of ad 7x 108 ves.," \ref{mod_pre} we show the evolution of the column densities of the observed species in a starless core (i.e. in the first scenario) with a density of $\times$ $^4$ and a FR=0.2, that corresponds to a percentage of CO on grains of at $\times$ $^6$ yrs."251" At lime around 9x 109 vrs the observed species stabilize their abundances and the predicted column densities of CS. and fall in the range of the observed values while NIL, and agree within a factor of 15."," At time around $\times$ $^6$ yrs the observed species stabilize their abundances and the predicted column densities of CS, and fall in the range of the observed values while $_3$ and agree within a factor of 15."252" A good agreement with the observations is also found al early limes. around 5.5x 109 vrs. for a slightly higher depletion efficiency (FR=0.4). however. also in this case. NIL, is between one and two orders of magnitude higher (hen observed."," A good agreement with the observations is also found at early times, around $\times$ $^6$ yrs, for a slightly higher depletion efficiency (FR=0.4), however, also in this case, $_3$ is between one and two orders of magnitude higher then observed."253 In Fig., In Fig.254 LO we show the column densities of the observed species vs time in (he best [fit model of the second scenario. ie. where a voung protostar is present in the core.," \ref{mod_pro} we show the column densities of the observed species vs time in the best fit model of the second scenario, i.e. where a young protostar is present in the core."255 The general behaviour of Cs. NII; and is similar to the previous scenario while and CIL;OLII behave differently.," The general behaviour of CS, $_3$ and is similar to the previous scenario while and $_3$ OH behave differently."256 In particular does not show the second peak and ΠΟΠ decreases quickly., In particular does not show the second peak and $_3$ OH decreases quickly.257 The reason why methanol, The reason why methanol258Galactic winds and outllows are the primary mechanism by which galaxies deposit energy. anc metalenrichecl gas into the intergalactic. medium. This can greatly. alfect. the evolution of the LGAL and the subsequent formation of other generations of galaxies.,"Galactic winds and outflows are the primary mechanism by which galaxies deposit energy and metal-enriched gas into the intergalactic medium This can greatly affect the evolution of the IGM, and the subsequent formation of other generations of galaxies."259 Feedback by galactic outflows can provide an explanation for the observed high mass-to-light ratio of dwarl galaxies ancl the abundance of dwarf galaxies in the Local Group. and can solve various problems with galaxy. formation mocels.," Feedback by galactic outflows can provide an explanation for the observed high mass-to-light ratio of dwarf galaxies and the abundance of dwarf galaxies in the Local Group, and can solve various problems with galaxy formation models,"260are significant: we find tvpical changes of oover (he full range of coronal temperatures for O abundance variations of a factor of 2.,are significant: we find typical changes of over the full range of coronal temperatures for O abundance variations of a factor of 2.261 When the other abundant light elements C ancl N are allowed to scale with the ο) abundance the diagnostic fares slightly less well but does provide discrimination between the photospherie compositions of GS and (2005)., When the other abundant light elements C and N are allowed to scale with the O abundance the diagnostic fares slightly less well but does provide discrimination between the photospheric compositions of GS and .262. The O abundance is llower than that of GS. and the ο Ka EWs differ by20-25%.," The O abundance is lower than that of GS, and the O $\alpha$ EWs differ by."263.. For a given exciting X-ray spectrum. changes in [Inorescent line strength with different photospheric parameters depend on changes in theshell of the fluorescent line in question.," For a given exciting X-ray spectrum, changes in fluorescent line strength with different photospheric parameters depend on changes in the of the fluorescent line in question."264 For a solar composition. two other sources of opacity in (the vicinitw of the O edge are C and N. The chemical compositions of GS and differ by «10 iin O/C and O/N ratios. ancl the lockstep changes in these elements that dilutes the ellect of the dillerent O abundances relative to I1 on the O ἵνα EW.," For a solar composition, two other sources of opacity in the vicinity of the O edge are C and N. The chemical compositions of GS and differ by $< 10$ in O/C and O/N ratios, and the lockstep changes in these elements that dilutes the effect of the different O abundances relative to H on the O $\alpha$ EW."265 Despite this slightly lower sensitivity of the [Inorescent line to the elobal chemical composition. accurate measurements of the O Ίνα EWs still potentially provide a new and relatively direct means of assessing the veracity of the GS and mixtures.," Despite this slightly lower sensitivity of the fluorescent line to the global chemical composition, accurate measurements of the O $\alpha$ EWs still potentially provide a new and relatively direct means of assessing the veracity of the GS and mixtures."266 The line EW is also sensitive to (he abundances adopted [or the exciting coronal spectrum., The line EW is also sensitive to the abundances adopted for the exciting coronal spectrum.267 This is due to the large contribution of the O VII He-lke complex to the source of ionising photons for coronal temperatures logZ'«6.5., This is due to the large contribution of the O VII He-like complex to the source of ionising photons for coronal temperatures $\log T < 6.5$.268 To a lesser extent. Fe L-shell ancl Ne. Meg and $i [-like and Ile-like lines also make a contribution for temperatures up to lopLT7.0.," To a lesser extent, Fe L-shell and Ne, Mg and Si H-like and He-like lines also make a contribution for temperatures up to $\log T\sim 7.0$."269 We have examüned (he sensitivity (o coronal abundances by comparison of O Ίνα line EWs computed for coronal spectra generated using GS and compositions. as illustrated in Figure 3..," We have examined the sensitivity to coronal abundances by comparison of O $\alpha$ line EWs computed for coronal spectra generated using GS and compositions, as illustrated in Figure \ref{f:ew}."270 The former are hieher than the latter by an amount that decreases [from [for temperatures logT<6.3 where the O VII lines dominate. to ~LOY aat logο6.8 1.0.," The former are higher than the latter by an amount that decreases from for temperatures $\log T \leq 6.3$ where the O VII lines dominate, to $\sim 10$ at $\log T \sim 6.8$ –7.0."271" Differences at higher temperatures are largely due to the lower Ne. Mg. Si and Fe abundances in the composition (bv 74. 12. 10 and respectively),"," Differences at higher temperatures are largely due to the lower Ne, Mg, Si and Fe abundances in the composition (by 74, 12, 10 and respectively)."272 In addition to uncertainties in the solar O content. coronal abundance variations are also expected as a result of chemical fractionation. in which the abundances of elements with low first lonisation potentials are seen to differ from photospheric values by. factors of up to ~11992)..," In addition to uncertainties in the solar O content, coronal abundance variations are also expected as a result of chemical fractionation, in which the abundances of elements with low first ionisation potentials are seen to differ from photospheric values by factors of up to $\sim2734$."274 In this context. we emphasise (hat in a practical application of the fIuorescence technique the photospheric abundance would be deduced by comparing the model fluorescent EWs computed for different photospheric abundances and (he," In this context, we emphasise that in a practical application of the fluorescence technique the photospheric abundance would be deduced by comparing the model fluorescent EWs computed for different photospheric abundances and the"27530% of the currently observed Li or 11erely be a secondary source of Li in the galaxy.,$\%$ of the currently observed Li or merely be a secondary source of Li in the galaxy.276the proximity of these lines. the best. model describes the entire observed spectrum (except Ho) very well.,"the proximity of these lines, the best model describes the entire observed spectrum (except $\alpha$ ) very well."277 This velocity measurement method. have multiple sources of error., This velocity measurement method have multiple sources of error.278 One of them may be the systematic bias due the approximations in the model (LII. power-LIaw atmosphere. simple source function. ete.).," One of them may be the systematic bias due the approximations in the model (LTE, power-law atmosphere, simple source function, etc.)."279 However. the comparison of our results with those [rom full NULL models (822)) show no systematic bias in the case of SNe 1999em and 2005ces.," However, the comparison of our results with those from full NLTE models \ref{sec_results}) ) show no systematic bias in the case of SNe 1999em and 2005cs."280 The agreement between the velocities from these two very cillerent modeling codes are within X10 percent., The agreement between the velocities from these two very different modeling codes are within $\pm 10$ percent.281 For SN 2006bp the dillerences are higher. but it will be shown below that for this SN the models do not describe well the spectral features we use. contrary tothe moclels (822)).," For SN 2006bp the differences are higher, but it will be shown below that for this SN the models do not describe well the spectral features we use, contrary tothe models \ref{sec_results06bp}) )."282 Another source of error may be the correlation between the parameters., Another source of error may be the correlation between the parameters.283 In Fig., In Fig.284 3. we present contour plots of the V hyperspace around its minimum. as a function of i14 and several other parameters that can allect the shape of the fitted A5169 feature.," \ref{contur} we present contour plots of the $\chi^2$ hyperspace around its minimum, as a function of $v_{model}$ and several other parameters that can affect the shape of the fitted $\lambda5169$ feature."285 The thick black contour curve corresponds to 50% higher V7 than the minimum value., The thick black contour curve corresponds to 50 higher $\chi^2$ than the minimum value.286" I is visible that correlation is indeed. present (1.0. the contours are distorted) between 0,,,5,4 ancl the power-law exponent n or the optical depth rp.", It is visible that correlation is indeed present (i.e. the contours are distorted) between $v_{model}$ and the power-law exponent $n$ or the optical depth $\tau_{Fe}$.287 The correlation is much less between Cyrede and Tp Of and Mg. whose features mav blend with A5169.," The correlation is much less between $v_{model}$ and $\tau_{ref}$ of and , whose features may blend with $\lambda5169$."288" Lowever. even for the correlated parameters. selecting n or τε, very far from their optimum value can alter 0,54 only bv a few hundred km +."," However, even for the correlated parameters, selecting $n$ or $\tau_{Fe}$ very far from their optimum value can alter $v_{model}$ only by a few hundred km $^{-1}$."289" Thus. we conclude that uncertainties in finding the minimum of 47 do not cause errors in co, that significantly exceed. the uncertainty due to the spectral resolution of the observed spectra (which is usually between 200 - 300 kim ly "," Thus, we conclude that uncertainties in finding the minimum of $\chi^2$ do not cause errors in $v_{model}$ that significantly exceed the uncertainty due to the spectral resolution of the observed spectra (which is usually between 200 - 300 km $^{-1}$ )."290A possible source of uncertainty may be that during the final fitting the wavelength interval around the used spectral feature is chosen somewhat subjectively., A possible source of uncertainty may be that during the final fitting the wavelength interval around the used spectral feature is chosen somewhat subjectively.291 However. our tests showed that changing the limits reasonably has negligible ellect on the final velocities.," However, our tests showed that changing the limits reasonably has negligible effect on the final velocities."292 ]t is emphasized. that. although the final fitting is restricted to a vicinity of a well-defined. spectral line. this method is certainly more reliable than the measurement of only the location ofthe οἱ the same feature.," It is emphasized that although the final fitting is restricted to a vicinity of a well-defined spectral line, this method is certainly more reliable than the measurement of only the location of the of the same feature."293 As it was discussed above. the minimum can be significantly and systematically altered by signal-to-noise. spectral resolution. blending. ete.," As it was discussed above, the minimum can be significantly and systematically altered by signal-to-noise, spectral resolution, blending, etc."294 The fitting of a model spectrum to the feature is expected to overcome these dillieulties. provided the underlving model is not too [ar from reality.," The fitting of a model spectrum to the feature is expected to overcome these difficulties, provided the underlying model is not too far from reality."295" UsingSYNOW as described. above. we determined. the best-fitting. parameters of all SNe spectra from See. οὃν,"," Using as described above, we determined the best-fitting parameters of all SNe spectra from Sec. \ref{sec_data}."296" The resulting mioclel velocities are collected in Table 11 in Appendix D. The best-fittingSYNOW parameters. such as T, For each atom/ion. the power-law exponent n and eycde together with the chosen Z,55,. can be found in Table in Appendix €. In Table Bl we also list theey, and Dus velocities."," The resulting model velocities are collected in Table \ref{vel} in Appendix B. The best-fitting parameters, such as $\tau_{ref}$ for each atom/ion, the power-law exponent $n$ and $v_{model}$ together with the chosen $T_{phot}$, can be found in Table \ref{synowtable} in Appendix C. In Table \ref{vel} we also list the$v_{Fe}$ and $v_{H\beta}$ velocities."297 For SNe 1999em. 2005es ancl 2006bp.. we collected. the photospheric velocities from: CMFGEN. models of Dessart&LHillier(2006) ancl Dessartetal.(2008).," For SNe 1999em, 2005cs and 2006bp, we collected the photospheric velocities from $\tt CMFGEN$ models of \citet{dessart2006} and \citet{dessart2008}."298. These are included in Table Bl as Όρη., These are included in Table \ref{vel} as $v_{nlte}$.299 Velocities from the ecross-correlation technique (Sect. 22)), Velocities from the cross-correlation technique (Sect. \ref{sec_cross}) )300 were obtained using two sets of template spectra., were obtained using two sets of template spectra.301 The first set contained the 22 observed spectra of SN. 1999em. (set. #11). while the second set was based on the CMFGEN mocels mentioned above (set 22).," The first set contained the 22 observed spectra of SN 1999em (set 1), while the second set was based on the $\tt CMFGEN$ models mentioned above (set 2)."302" The velocities ofthe template spectra were re, for set #11 and 0,5, for set #22.", The velocities ofthe template spectra were $v_{Fe}$ for set 1 and $v_{nlte}$ for set 2.303 We cross-correlated all the observed spectra with the two sets separately on the wavelength range of 4500 5500Α.. and the resulting velocities are also in Table BI as v.," We cross-correlated all the observed spectra with the two sets separately on the wavelength range of 4500 – 5500, and the resulting velocities are also in Table \ref{vel} as $v_{cc}$."304 Fie., Fig.305 4 shows tna against phase for all studied SNe (top left. panel). and the ratio of tia to all the other velocities.," \ref{velocities} shows $v_{model}$ against phase for all studied SNe (top left panel), and the ratio of $v_{model}$ to all the other velocities."306 The caleulated: velocities all show the expected decline with phase as the photosphere moves deeper anc deeper within the ejectra. toward slower expanding lavers.," The calculated velocities all show the expected decline with phase as the photosphere moves deeper and deeper within the ejectra, toward slower expanding layers."307 Similar plots containing the ratio tyronefle. and Robsflee as Cunetions of phase. are presented in Fig. 5..," Similar plots containing the ratio $v_{model} / v_{cc}$ and $v_{abs} / v_{cc}$ as functions of phase, are presented in Fig. \ref{velcc}."308 In the followings we provide some details of deriving these velocities for cach object ancl discuss some object-specilie dillerences between them., In the followings we provide some details of deriving these velocities for each object and discuss some object-specific differences between them.309" When determining t,o. Wilh SYNOW.. £2 was fitted or the first 6 spectra. then the A5169 feature was used or the remaining 16 spectra."," When determining $v_{model}$ with , $H\beta$ was fitted for the first 6 spectra, then the $\lambda5169$ feature was used for the remaining 16 spectra."310 The resulting velocities are tween L1000 and 1800 kms 1., The resulting velocities are between $11000$ and $1800$ km $^{-1}$.311 As seen in the bottom right xiuiel of Eig. 4.. ," As seen in the bottom right panel of Fig. \ref{velocities}, ,"312μι and μι are about the same for the carly phases (before the appearance of the lines). while ater cus tends to be higher than Όροι.," $v_{model}$ and $v_{H\beta}$ are about the same for the early phases (before the appearance of the lines), while later $v_{H\beta}$ tends to be higher than $v_{model}$."313" Also. between day [15 and dày |40. 60,27 ds à slightly higher than er, (Fig."," Also, between day +15 and day +40, $v_{model}$ is a slightly higher than $v_{Fe}$ (Fig."314 4 bottom left panel)., \ref{velocities} bottom left panel).315" After day |40 6,5 drops below ey, and their ratio Increases toward later phases.", After day +40 $v_{model}$ drops below $v_{Fe}$ and their ratio increases toward later phases.316 The velocities fron: models of (Dessart&Llill, The velocities from models of \citep{dessart2006} (Fig.317"ier2006) (Lig. 4 top right panel) agree with ey,cuc7.", \ref{velocities} top right panel) agree with $v_{model}$.318 Phe cross-Correlation with set. #22 (Fie.5 bottom: panels) gave similar results for the first few points. but. overestimate Cone between davs |22 and. |80.," The cross-correlation with set 2 \ref{velcc} bottom panels) gave similar results for the first few points, but overestimate $v_{model}$ between days +22 and +80."319" ""ον mostly fall between Cg sand epo which is expected. since we eross-correlated the range of 4500 0500A.. where these features appear."," They mostly fall between $v_{H\beta}$ and $v_{Fe}$, which is expected, since we cross-correlated the range of 4500 – 5500, where these features appear."320 The SYNOW model velocities of the L1 spectra that cover the second of halfthe plateau phase are between  3400 and 1700 km s+., The model velocities of the 11 spectra that cover the second half of the plateau phase are between $\sim$ 3400 and 1700 km $^{-1}$.321 These are similar to those of SN 1999em at the same phase., These are similar to those of SN 1999em at the same phase.322 Both ep. and eg; ave higher thaney. at all epochs. especially the latter with a factor of about 1.8 (Fig. 4)).," Both $v_{Fe}$ and $v_{H\beta}$ are higher than$v_{model}$ at all epochs, especially the latter with a factor of about $1.8$ (Fig. \ref{velocities}) )."323 No model was available for SN. 2004dj., No model was available for SN 2004dj.324 C'ross- with both template sets gavevery similar results., Cross-correlation with both template sets gavevery similar results.325" They are only slightly higher than both ruc. and vj, (Fig.5)).", They are only slightly higher than both $v_{model}$ and $v_{Fe}$ \ref{velcc}) ).326 Forthe first 6 spectra the model was optimized for £13. then for the A5169 feature.," Forthe first 6 spectra the model was optimized for $H\beta$ , then for the $\lambda 5169$ feature."327 Theresulting model velocities ave between 9700 and 1800 kms + (Fig. 4))., Theresulting model velocities are between 9700 and 1800 km $^{-1}$ (Fig. \ref{velocities}) ).328 The, The329of GALEX | 0117 by comparing our non-L'TIS analysis to LEE results (Section 3.3.2).,of GALEX $+$ 0117 by comparing our non-LTE analysis to LTE results (Section 3.3.2).330 Our original spectral svntheses were computed with the ine list available on the CD-ROAL No., Our original spectral syntheses were computed with the line list available on the CD-ROM No.331 23 of Ixurucz&Bell(1995) /kurucz23/sekur., 23 of \citet{kur1995} .332html.. For each ton. we now compare he oscillator strengths. (fi) and Stark line. broadening xwameters. (D) listed. by Ixurucz&Bell(1995). to the rest available theoretical and. experimental data.," For each ion, we now compare the oscillator strengths $f_{\rm ij}$ ) and Stark line broadening parameters $\Gamma$ ) listed by \citet{kur1995} to the best available theoretical and experimental data."333" Data on ine oscillator strengths are also available at the National Institute of Standards and Lechnoloey We noted that the isotopic shift in the 7°21 Πλοτος doublet is only [0.85 (Drullinger.Wineland.&Jerequist1980). and. therefore. we do not expect observable ellects on the MgiAHSI doublet in ""resolution spectra."," Data on line oscillator strengths are also available at the National Institute of Standards and Technology We noted that the isotopic shift in the $^{26-24}$ $\lambda$ 2798 doublet is only $+0.85$ $^{-1}$ \citep{dru1980}334 and, therefore, we do not expect observable effects on the $\lambda$ 4481 doublet in $^{-1}$ -resolution spectra."335 Similarly. other isotopic shifts (ce... οοἱ 781) may be neglected in the present study. (seeBerengut.Dzuba.&Flambaum 2003)..," Similarly, other isotopic shifts (e.g., $^{30}{\rm Si}/^{28}{\rm Si}$ ) may be neglected in the present study \citep[see][]{ber2003}. ."336 The rich silicon. lino. spectrum in GALEN J1931[0117 prompted a detailed review of available data., The rich silicon line spectrum in GALEX J1931+0117 prompted a detailed review of available data.337 Table 2. lists and compares fi from popular data compilations (CD23 and NIST) to homogeneous. theoretical. or experimental data sets., Table \ref{tbl-2} lists and compares $f_{\rm ij}$ from popular data compilations (CD23 and NIST) to homogeneous theoretical or experimental data sets.338.. We. noted considerable variations in oscillator strengths. in particular in the À3862 triplet and in A5041.024.," We noted considerable variations in oscillator strengths, in particular in the $\lambda$ 3862 triplet and in $\lambda$ 5041.024."339. Discrepancies of the order of 40 to would. allect individual abundance measurements in equal measures., Discrepancies of the order of 40 to would affect individual abundance measurements in equal measures.340 For example. the ratio of CD23 data to experimental fi; values is 1.06 but varies with a standard deviation σ—38'4 while the NIST data that are largely based on the experimental values vary by only with an average ratio of 0.98. and the theoretical data (Artruetal.LOST) vary by with an average ratio of 0.98.," For example, the ratio of CD23 data to experimental $f_{\rm ij}$ values is 1.06 but varies with a standard deviation $\sigma=$ while the NIST data that are largely based on the experimental values vary by only with an average ratio of 0.98, and the theoretical data \citep{art1981} vary by with an average ratio of 0.98."341 The adoption of one data set over another will not have a large effect on the average abundance. but. individual line measurements are less reliable.," The adoption of one data set over another will not have a large effect on the average abundance, but individual line measurements are less reliable."342 Strong saturated lines are sensitive to line broadening parameters., Strong saturated lines are sensitive to line broadening parameters.343 “Table 3/— lists full-width. at. half-maximunm: (FWHAI= Pu) Stark widths and shifts due to electron impacts for strong optical lines., Table \ref{tbl-3} lists full-width at half-maximum $\equiv 2w$ ) Stark widths and shifts due to electron impacts for strong optical lines.344" Ixurucz&Bell(1995) tabulate the circular. frequeney per electron. P that we converted into the PWLIM at 5,=107 ? using the formula: where e is the speed of light.", \citet{kur1995} tabulate the circular frequency per electron $\Gamma$ that we converted into the FWHM at $n_e=10^{17}$ $^{-3}$ using the formula: where $c$ is the speed of light.345Lanz.Dimitrijevic.&Artru(1988) tabulates FWILAT values for electron. and. proton impacts separately., \citet{lan1988} tabulates FWHM values for electron and proton impacts separately.346 For most lines the proton contribution to the total width is 21054 up to1, For most lines the proton contribution to the total width is $\approx$ up to.3475/4... ENIM from ]|xurucz&Bell(1995). are on average Thelareer than the experimental values with a standard deviation of while the values from Lanz.Dimitrijevic.&Artru(1955) are only lower than the experimental values with a stanelare deviation of21%., The FWHM from \citet{kur1995} are on average larger than the experimental values with a standard deviation of while the values from \citet{lan1988} are only lower than the experimental values with a standard deviation of.348 The execllent agreement between Lanz.Dimitrijevic.&Artru(1988). and the experiments prompted. us to explore two options in computing detailed. silicon line spectra., The excellent agreement between \citet{lan1988} and the experiments prompted us to explore two options in computing detailed silicon line spectra.349 In option Lowe adopted. the silicon oscillator strengths. from Artruetal.(1981) and the line broadening parameters for electrons anclprotons from Lanz.Dimitrijevic.&Artru(1988)., In option 1 we adopted the silicon oscillator strengths from \cite{art1981} and the line broadening parameters for electrons andprotons from \citet{lan1988}.350". The aepening parameters are tabulated at 5. 10. 20. and C dx with estimated: uncertainties of less than at 5»,=107 em (or a depth zj£z2 in the atmosphere)."," The line broadening parameters are tabulated at 5, 10, 20, and $\times10^3$ K with estimated uncertainties of less than at $n_e=10^{17}$ $^{-3}$ (or a depth $\tau_R\approx 2$ in the atmosphere)."351" Also. the clleet of Stark shifts are. included using the experimental data of Gonzálezctal.(2002) although we have no information onthe scaling of d, with temperature or on M magnitude of Stark shifts due to ions (protonsin this case)."," Also, the effect of Stark shifts are included using the experimental data of \citet{gon2002} although we have no information onthe scaling of $d_e$ with temperature or on the magnitude of Stark shifts due to ions (protons in this case)."352 The ellect of Stark shifts is apparent in radial velocity measurements of the AADOVAL-5055 ancl AA5967-5978 MNopts (seeTable1 ., The effect of Stark shifts is apparent in radial velocity measurements of the $\lambda\lambda$ 5041-5055 and $\lambda\lambda$ 5967-5978 multiplets \citep[see Table~\ref{tbl-1} .353 byIn option 2. we emploved the cata Cfi; and D) iil lxurucz (1995)... and neglected the ellect of Stark shifts.," In option 2, we employed the data $f_{\rm ij}$ and $\Gamma$ ) compiled by \citet{kur1995}, and neglected the effect of Stark shifts."354 Unlike silicon. the magnesium abundance measurement. is based on a single doublet.," Unlike silicon, the magnesium abundance measurement is based on a single doublet."355 The red. Mg11AT989. multiplet shows evidence of Stark shifts (Vennes.Ixawka.&Németh2010a) although the spectrum is particularly noisy in the vicinity. of the multiplet.," The red $\lambda$ 7989 multiplet shows evidence of Stark shifts \citep{ven2010a}356 although the spectrum is particularly noisy in the vicinity of the multiplet."357 On the other hand. the strong TAKS1 doublet is well exposed.," On the other hand, the strong $\lambda$ 4481 doublet is well exposed."358 Because of its strength. the Dine profile is particularly sensitive to. broadening parameters. although quoted oscillator strength: values.are consistent (NISTandIxurucz&Bell1995) within = 2%..," Because of its strength, the line profile is particularly sensitive to broadening parameters, although quoted oscillator strength valuesare consistent \citep[NIST and][]{kur1995}359 within $\la2$ ."360 The compilation of Ixurucz&Bell) does not provide a Stark width for the Mgl1A4481. doublet. which is then estimated in using the formula (Castelli2005) ," The compilation of \citet{kur1995} does not provide a Stark width for the $\lambda$ 4481 doublet which is then estimated in using the formula \citep{cas2005}361 "362applicable to large gas-grain chemical networks.,applicable to large gas-grain chemical networks.363 Such a scheme must provide rate continuity over the stochastic-deterministic threshold at (N(7))=1., Such a scheme must provide rate continuity over the stochastic--deterministic threshold at $\langle N(i) \rangle = 1$.364 Any functional form must also be integrable using standard differential equation-solving techniques (i.e. the Gear algorithm)., Any functional form must also be integrable using standard differential equation-solving techniques (i.e. the Gear algorithm).365" Below, further modifications to, and restrictions on, the basic equations (11) and (12) are formulated."," Below, further modifications to, and restrictions on, the basic equations (11) and (12) are formulated."366 The switch-over between the modified rate and the standard rate may produce a discontinuity in the rates at (N(i))=1., The switch-over between the modified rate and the standard rate may produce a discontinuity in the rates at $\langle N(i) \rangle =1$.367" Dependent on the relative rates of all the processes involved, this may present an impediment to accurate calculations."," Dependent on the relative rates of all the processes involved, this may present an impediment to accurate calculations."368" Therefore, a simple empirical function, f, is introduced to make the transition smoother whilst preserving a fast switch-over."," Therefore, a simple empirical function, $f$, is introduced to make the transition smoother whilst preserving a fast switch-over."369" Under this scheme, production rates arealways modified according to: where: This has the effect that when (N(A)),(N(B))«1, the rate is always ""stochastic"", whilst quickly tending towards the deterministic rate as either (N(A)) or (N(B)) rises above unity."," Under this scheme, production rates are modified according to: where: This has the effect that when $\langle N(A) \rangle, \langle N(B) \rangle < 1$, the rate is always “stochastic”, whilst quickly tending towards the deterministic rate as either $\langle N(A) \rangle$ or $\langle N(B) \rangle$ rises above unity."370" For example, if species A and B both attain abundances of 10 atoms/molecules per grain, the deterministic contribution to the total production rate will be of the normal deterministic rate."," For example, if species $A$ and $B$ both attain abundances of 10 atoms/molecules per grain, the deterministic contribution to the total production rate will be of the normal deterministic rate."371 Expressions (14) (15) therefore allow the modified rate contribution to replace a fraction of the total deterministic rate that corresponds to the reaction of 1 atom/molecule of species A with 1 atom/molecule of species B., Expressions (14) (15) therefore allow the modified rate contribution to replace a fraction of the total deterministic rate that corresponds to the reaction of 1 atom/molecule of species $A$ with 1 atom/molecule of species $B$.372" Equation (13) is also applied, meaning that if Rjo4(AB)>kap:(N(A))-(NCB)), the total rate is equal to the unmodified deterministic rate."," Equation (13) is also applied, meaning that if $R_{mod}(AB) > k_{AB} \cdot \langle N(A) \rangle \cdot \langle N(B) \rangle$, the total rate is equal to the unmodified deterministic rate."373" The continuous rate-modification scheme outlined above is used to examine the water and methanol systems, at 15 K, as investigated by Barzel&Biham(2007b)."," The continuous rate-modification scheme outlined above is used to examine the water and methanol systems, at 15 K, as investigated by \cite{barzel2}."374". The former scheme includes reactions between surface species H, O, and OH, resulting in Hz, O» and H2O production."," The former scheme includes reactions between surface species H, O, and OH, resulting in $_2$, $_2$ and $_2$ O production."375" The latter scheme also includes reactions with CO, HCO, H2CO, and CH30, leading to the production of methanol, CH30H, and carbon dioxide, CO»."," The latter scheme also includes reactions with CO, HCO, $_2$ CO, and $_3$ O, leading to the production of methanol, $_3$ OH, and carbon dioxide, $_2$."376" The reactions, fluxes, and binding energies are indicated in Tables 1 2."," The reactions, fluxes, and binding energies are indicated in Tables 1 2."377 It should be noted that the methanol system investigated by Barzel Biham does not include any activation-energy barriers., It should be noted that the methanol system investigated by Barzel Biham does not include any activation-energy barriers.378" In fact, activation energies substantially complicate the behaviour of the methanol system; see Section 6."," In fact, activation energies substantially complicate the behaviour of the methanol system; see Section 6."379" Figures 3 4 show population sizes and production rates of key species, for the water and methanol systems, respectively."," Figures 3 4 show population sizes and production rates of key species, for the water and methanol systems, respectively."380 Populations obtained from the new method (Section 4.1) are very well matched to the master-equation results of Barzel Biham., Populations obtained from the new method (Section 4.1) are very well matched to the master-equation results of Barzel Biham.381" Production rates are accurate at high and very low values of S, but values near the stochastic-deterministic threshold are more obviously inaccurate, whilst obeying the correct trend."," Production rates are accurate at high and very low values of $S$, but values near the stochastic–deterministic threshold are more obviously inaccurate, whilst obeying the correct trend."382" However, the new method is a very good first approximation."," However, the new method is a very good first approximation."383 What are the underlying physical reasons for the disagreements?, What are the underlying physical reasons for the disagreements?384 The production rates obtained from the new method rise to the rate equation values at lower values of S than the master-equation results., The production rates obtained from the new method rise to the rate equation values at lower values of $S$ than the master-equation results.385" This occurs before any reactants approach a population of 1, as the rate limit of equation (13) is reached before this."," This occurs before any reactants approach a population of 1, as the rate limit of equation (13) is reached before this."386" Hence, the empirical function that is used to switch over at the (N(i))=1 threshold is not the cause."," Hence, the empirical function that is used to switch over at the $\langle N(i) \rangle = 1$ threshold is not the cause."387" In fact, the modified rates are too fast because competition between surface processes has not been considered."," In fact, the modified rates are too fast because competition between surface processes has not been considered."388" At low values of S, reactions are extremely fast, due to the fast reaction rates of all reactants; see equation (3)."," At low values of $S$, reactions are extremely fast, due to the fast reaction rates of all reactants; see equation (3)."389" As S increases, reaction becomes slower, and the possibility arises that one or other reactant may evaporate before the two meet in the same binding site and react."," As $S$ increases, reaction becomes slower, and the possibility arises that one or other reactant may evaporate before the two meet in the same binding site and react."390" For hydrogen atoms in the water and methanol systems, using the binding energies shown in Table 2, Κηχ=Kevap(H) when S~460 sites per grain."," For hydrogen atoms in the water and methanol systems, using the binding energies shown in Table 2, $k_{HX}=k_{evap}(H)$ when $S \simeq 460$ sites per grain."391Our individual limits are siguificautlv better than the studies at 2=3. πο if there are significaut nunibers of ealaxies with lower escape fractious (eg. fiero)~0.20.,"Our individual limits are significantly better than the studies at $z=3$, so if there are significant numbers of galaxies with lower escape fractions (eg. $f_{\mathrm{esc,rel}}\sim 0.20$,"392 rather than unitv). we would be able to detect thei.," rather than unity), we would be able to detect them."393 We see no evidence for this scenario iu our sample., We see no evidence for this scenario in our sample.394 Comparing our findings. which are in agreement with all other 2<1 f; measurements. with :~23 studies (?7) nuplies that the average escape fraction evolves with redshift. but the cause of this evolution reais nuknown.," Comparing our findings, which are in agreement with all other $z<1$ $f_{\mathrm{esc}}$ measurements, with $z\sim3$ studies \citep{2006ApJ...651..688S,2009ApJ...692.1287I} implies that the average escape fraction evolves with redshift, but the cause of this evolution remains unknown."395 It should be noted. however. that foreground contanunation which is likely more severe at higher redshift. ταν account for some of the apparent evolution seen between 2= Laud 3. by this work aud others (??7)..," It should be noted, however, that foreground contamination which is likely more severe at higher redshift, may account for some of the apparent evolution seen between $z=1$ and 3, by this work and others \citep{2010arXiv1001.3412S,2010MNRAS.404.1672V,2006MNRAS.371L...1I}."396 We proceed by focusing on possible imechauisunis that could explain the lack of large escape fraction galaxies at lol., We proceed by focusing on possible mechanisms that could explain the lack of large escape fraction galaxies at $z\sim1$.397 When comparing biel and low-redshift ealaxy saluples. there is always some degree of uucertaimty reearding the true analog nature of the two populations.," When comparing high- and low-redshift galaxy samples, there is always some degree of uncertainty regarding the true analog nature of the two populations."398 As discussed in Section 2 and shown in Figure | we selected a sample of LBC analogs sharing many of he same properties of the 7. and ? 2—3 LBCs.," As discussed in Section \ref{sec:selection} and shown in Figure \ref{fig:surf_bright}, we selected a sample of LBG analogs sharing many of the same properties of the \citet{2006ApJ...651..688S} and \citet{2009ApJ...692.1287I} $z\sim3$ LBGs."399 Figure 10 further hiehliehts their similarities showiug he distribution of reddening. stellar mass and rest-παλιο UV luminosity in these sources is simular to the distribution in LBCs.," Figure \ref{fig:ebv_mass} further highlights their similarities showing the distribution of reddening, stellar mass and rest-frame UV luminosity in these sources is similar to the distribution in LBGs."400 The similarity in mass together with the UVoptical colors of the UVLCs miplies that hey iav still be uudergoiusg an carly. major episode of star formation rather than a small burst on top of a iddenu older population (alsosee?.forthesamereason LBGs}..," The similarity in mass together with the UV–optical colors of the UVLGs implies that they may still be undergoing an early, major episode of star formation rather than a small burst on top of a hidden older population \citep[also see][for the same reason applied401to LBGs]{2004ApJS..154...97B}."402" Ultimately, this sample of LBC analogs shares Muuerous smiluities to the parent LBC population. but since only —105415% of LBCs have been observed with sguificaut LyC detections. perhaps this subclass of LBCs has other processes at work allowing or aidius iu the escape of LyC plotous."," Ultimately, this sample of LBG analogs shares numerous similarities to the parent LBG population, but since only $\sim10\%-15\%$ of LBGs have been observed with significant LyC detections, perhaps this subclass of LBGs has other processes at work allowing or aiding in the escape of LyC photons."403 Galaxy niergers offer au intriguing explanation for the Increased escape fraction seen at 2~3., Galaxy mergers offer an intriguing explanation for the increased escape fraction seen at $z\sim3$.404 7. noted that UVLGs typically exhibit faint tidal features sugecstive of aimeregcr or recent interaction., \citet{2008ApJ...677...37O} noted that UVLGs typically exhibit faint tidal features suggestive of a merger or recent interaction.405 They therefore propose that the super starbursts in LDCs are trigeered by eas-rich mergers., They therefore propose that the super starbursts in LBGs are triggered by gas-rich mergers.406 Similarly. ?— showed that of LBCs have structures akin to local starburst merecrs aud may be driven by similar processes.," Similarly, \citet{2009AJ....138..362P} showed that of LBGs have structures akin to local starburst mergers and may be driven by similar processes."407 As ealaxies collide. strong eravitational auc tidal forces can expel loug streams of stars. and ignite violent starbursts at rates of a few to hundreds of M. vrt (?77)..," As galaxies collide, strong gravitational and tidal forces can expel long streams of stars, and ignite violent starbursts at rates of a few to hundreds of $M_{\odot}$ $^{-1}$ \citep{1982ApJ...252..455S,2000ApJ...530..660B,2010ApJ...709.1067B}."408 During a merger. the tidal fields distort the ealaxies racially. drawing out galactic material iuto long tails. pluues aud bridges (e.g...2?)..," During a merger, the tidal fields distort the galaxies radially, drawing out galactic material into long tails, plumes and bridges \citep[e.g.,][]{1972ApJ...178..623T,1996ApJ...464..641M}."409 The ID I reservoirs can become disturbed. aud the neutral gas pulled away from the sources of ioniziug radiation producing low-cohuun deusitv lines of sight (??) through the ealaxics. in turn allowing the escape of LyC photons.," The H I reservoirs can become disturbed, and the neutral gas pulled away from the sources of ionizing radiation producing low-column density lines of sight \citep{2000AJ....119.1130H,2008AJ....135..548D} through the galaxies, in turn allowing the escape of LyC photons."410 Simulations bv? sugeest that the escape fraction in major merecrs can be large (οντως 30%) compared to nonduergers Gf< 10%) alone specific lines of sight.," Simulations by \citet{2008ApJ...672..765G} suggest that the escape fraction in major mergers can be large $f_{\mathrm{esc,rel}}\gsim30\%$ ) compared to non-mergers $f_{\mathrm{esc,rel}}<10\%$ ) along specific lines of sight."411 Within our suuple of 32 sealaxies. 1l had morphologies consistent with merecr activity.," Within our sample of 32 galaxies, 11 had morphologies consistent with merger activity."412 We independently stacked cight of these spectra (removing three due to their larger spatial extent) aud fud fire<2% (80 upper linüt).," We independently stacked eight of these spectra (removing three due to their larger spatial extent) and find $f_{\mathrm{esc,rel}}<2\%$ $3\sigma$ upper limit)."413 Tf mereiue is a viable mechanisin for clearing pathways in the ISM for LyC photons. the orientation of the svstem along the line of sight is also a likely factor. requiring a laree suple of UV. luminous mergers," If merging is a viable mechanism for clearing pathways in the ISM for LyC photons, the orientation of the system along the line of sight is also a likely factor, requiring a large sample of UV luminous mergers."414 Therefore. we cannot sav whether mergers are an important factor as our saple size is at present too sinall.," Therefore, we cannot say whether mergers are an important factor as our sample size is at present too small."415 Currently. there is a lack of deep high-resolution rest-frame optical nuaeine ofthe LDGs with larger escape fractions. aud the interpretation of UV morphologies remains problematic (?)..," Currently, there is a lack of deep high-resolution rest-frame optical imaging of the LBGs with larger escape fractions, and the interpretation of UV morphologies remains problematic \citep{2007ApJ...656....1L}."416 Future near-IR observations with oof the :~3 LBC leakers will shed light ou this hypothesis., Future near-IR observations with of the $z\sim3$ LBG leakers will shed light on this hypothesis.417 As discussed carlicr in the section. galaxy mergers are capable of clearing patlwavs. exposing UV bright stars.," As discussed earlier in the section, galaxy mergers are capable of clearing pathways, exposing UV bright stars."418 If merecrs do facilitate the escape of LyC radiation heu an evolving merger rate. may be responsible for he observed evolution in [οταν ," If mergers do facilitate the escape of LyC radiation then an evolving merger rate, may be responsible for the observed evolution in $f_{\mathrm {sc,rel}}$."419Numerous observational studies auc simulations haven shown that the ealaxv nerecr rate evolves with+ redshift.a eoiug. as ~(1+|:)27H CQUNTTTTy.," Numerous observational studies and simulations haven shown that the galaxy merger rate evolves with redshift, going as $\sim(1+z)^{2-3}$ \citep{2001ApJ...546..223G,2003AJ....126.1183C,2007ApJS..172..320K,2007ApJ...659..976H,2007ApJ...659..931B,2008MNRAS.386..909C,2010ApJ...709.1067B}."420 The factor of 3-1 iuerease m iuerger rate jetween id and 3 as seen observationally. would also increase the number of lines of sights aud range of encounter paramcters observed in 2—3 galaxy mergers o» the same factor.," The factor of 3-4 increase in merger rate between $z\sim$ 1 and 3 as seen observationally, would also increase the number of lines of sights and range of encounter parameters observed in $z\sim3$ galaxy mergers by the same factor."421 This would in tur increase the ikelihood of detecting LyC at higher redshift., This would in turn increase the likelihood of detecting LyC at higher redshift.422 The LyC€ escape fraction is limited by the distribution of neutral lvdrogen along a Lue of sight aud likely depends on galactic parameters;, The LyC escape fraction is limited by the distribution of neutral hydrogen along a line of sight and likely depends on galactic parameters.423 We now consider what ealaxy properties could evolve with redshift that reduce he cficiency of galactic outflows/climmeys in leaking LvC€ photons from Iuuinous galaxies., We now consider what galaxy properties could evolve with redshift that reduce the efficiency of galactic outflows/chimneys in leaking LyC photons from luminous galaxies.424 Typical galaxies Gucluding UW bright galaxies) have )en Shown to be 1.5-3 times smaller at 2~3 than heir local counterparts (2??)..," Typical galaxies (including UV bright galaxies) have been shown to be 1.5-3 times smaller at $z\sim3$ than their local counterparts \citep{2006ApJ...650...18T,2002ApJ...579L...1P,2004ApJ...600L.107F}."425 Although our sample was selected to have similar UV surface brightuesses as 2~3 LBCs (refer to Figure 1)). little is known about the rue optical sizes of LBGs or the gas distribution.," Although our sample was selected to have similar UV surface brightnesses as $z\sim3$ LBGs (refer to Figure \ref{fig:surf_bright}) ), little is known about the true optical sizes of LBGs or the gas distribution."426 The velocities of galactic winds or outilows have been fond to © proportional to the SFR in LBCs (?).. therefore LBC and LBG analogs. having similar SER should ium principle eenerate outflows with similar velocities (a few huudred Sos ly," The velocities of galactic winds or outflows have been found to be proportional to the SFR in LBGs \citep{2006MNRAS.373..571F}, therefore LBG and LBG analogs, having similar SFR should in principle generate outflows with similar velocities (a few hundred km $^{-1}$ )."427" However. smaller salaxies would have higher SERs per uuit volume. which cau result iu more efficient ealactic winds (7). more casily clearing pathways or ""ehinmnev-like"" structures. aud im turni allowing for higher ""eaol (?).."," However, smaller galaxies would have higher SFRs per unit volume, which can result in more efficient galactic winds \citep{2005ARA&A..43..769V}, more easily clearing pathways or “chimney-like” structures, and in turn allowing for higher $f_{\mathrm{esc,rel}}$ \citep{2003ApJ...599...50F}."428 With smaller galaxies. aud ligher-deusity starbursts comes the potential for a larger fraction of stars born iu very colpact star clusters. including super star clusters (SSCs).," With smaller galaxies, and higher-density starbursts comes the potential for a larger fraction of stars born in very compact star clusters, including super star clusters (SSCs)."429 SSCs cau have thousands of vouung (50M) stars within a μαΠο radius of —10pc (7)., SSCs can have thousands of young $<$ 50Myr) stars within a half-light radius of $\sim$ 10pc \citep{2005ARA&A..43..769V}.430 These extreme concentrations of hot O aud D stars can ercatly Hupact the state of the ISAL driving powerful galactic winds (like those seen iu M82). openime chanucls for LvC photons to escape.," These extreme concentrations of hot O and B stars can greatly impact the state of the ISM driving powerful galactic winds (like those seen in M82), opening channels for LyC photons to escape."431 SSCs have been detected in the tidal tails (2). and outer regions of galaxies. which could explain the spatially offset LyC cussion (to the optical Cluission) detected bv ? ina few 2~3 LBCs.," SSCs have been detected in the tidal tails \citep{2009AAS...21334401C} and outer regions of galaxies, which could explain the spatially offset LyC emission (to the optical emission) detected by \citet{2009ApJ...692.1287I} in a few $z\sim3$ LBGs."432 There is also some evidence that SSCs found in the local group, There is also some evidence that SSCs found in the local group433SO galaxies under study relevant to this analysis. as follows: Co.,"S0 galaxies under study relevant to this analysis, as follows: Col."434 1: the galaxy. denomination: Co., 1: the galaxy denomination; Col.435 2: alternate (NGC/IC) galaxy. denomination: Col., 2: alternate (NGC/IC) galaxy denomination; Col.436 3: the morphological classification: Co., 3: the morphological classification; Col.437 4: the H-band etlective radius rg and its error: Col., 4: the H-band effective radius $r_{e H}$ and its error; Col.438 5: the classification according to the IH-band prolile «lecomposition (see Col., 5: the classification according to the H-band profile decomposition (see Col.439 5 of Tab., 5 of Tab.440 2): , 2); Col.4416: the 1I-band bulge-to-total luminosity ratio (sce Col., 6: the H-band bulge-to-total luminosity ratio $B/T_H$ (see Col.442 6 of Tab., 6 of Tab.443 2): ., 2); Col.444 T: the D-band elfective radius rg and its error: ., 7: the B-band effective radius $r_{e B}$ and its error; Col.445Sithe classification according to the D-band profile decomposition (see Col., 8: the classification according to the B-band profile decomposition (see Col.446 5 of Tab., 5 of Tab.447 2): , 2); Col.4489: the B-band bulge-to-total luminosity ratio p (see Col., 9: the B-band bulge-to-total luminosity ratio $B/T_B$ (see Col.449 6 of Tab., 6 of Tab.450 2): , 2); Col.45110: the observed. total L-banc magnitude LL and its error: Col., 10: the observed total H-band magnitude H and its error; Col.452 11: the observed total B-band magnitude D and its CDDOL., 11: the observed total B-band magnitude B and its error.453 llereafter. no correction. for Galactic extinction. in direction either of the Virgo cluster or of the Coma cluster will be applied to the B- and L-bancl magnitudes. since this correction is negligible for our purposes.," Hereafter no correction for Galactic extinction in direction either of the Virgo cluster or of the Coma cluster will be applied to the B- and H-band magnitudes, since this correction is negligible for our purposes."454 No correction [or internal extinction to the photometric parameters. of individual galaxies will be applied either., No correction for internal extinction to the photometric parameters of individual galaxies will be applied either.455 lig., Fig.456 2 shows the distribution of the LS VCC early-tvpe clwarls listed in Tab., 2 shows the distribution of the 18 VCC early-type dwarfs listed in Tab.457 1 in the plane defined by the decimal ogarithm of the ratio o£ reg and rg (regir ag) and by the otal color index LL. Individual galaxies are represented w empty circles. asterisks or filled. circles if their surface xightness profile follows. either a ce Vaucouleurs-Iaw. a | disk decomposition. or an exponential-Iaw (tvpe 1. 2 or 3. respectively cf," 1 in the plane defined by the decimal logarithm of the ratio of $r_{e B}$ and $r_{e H}$ $r_{e B}/r_{e H}$ ) and by the total color index $-$ H. Individual galaxies are represented by empty circles, asterisks or filled circles if their surface brightness profile follows either a de Vaucouleurs-law, a $+$ disk decomposition, or an exponential-law (type 1, 2 or 3, respectively – cf."458 Sect., Sect.459 2)., 2).460 In panels a and sh) objects are classified. according to profile decomposition either in he IHI-band or in the B-band (cf., In panels `a' and `b' objects are classified according to profile decomposition either in the H-band or in the B-band (cf.461 Tab., Tab.462 2). respectively.," 2), respectively."463 Fig., Fig.464 2 shows that:, 2 shows that:465The relation. between far-infrarecl emission. and. the racio emission from galaxies at all redshifts is surprisingly tight (deJongetal.1985:HelouConclon1992:Ciar-pett2002) and leads to the conclusion that both trace recent star-[ormation. in the local ancl distant Universe.,"The relation between far-infrared emission and the radio emission from galaxies at all redshifts is surprisingly tight \citep{deJong85, Helou85, Condon92, Garret02} and leads to the conclusion that both trace recent star-formation, in the local and distant Universe."466 ‘The far-infrared cmiussion is believed. to arise. from. the thermal emission of dusty clouds surrounding regions of star formation. whereas the radio emission arises [ron cosmic-ray electrons. accelerated in supernova remnants. of the dying stars. which emit svachrotron racliation.," The far-infrared emission is believed to arise from the thermal emission of dusty clouds surrounding regions of star formation, whereas the radio emission arises from cosmic-ray electrons accelerated in supernova remnants of the dying stars, which emit synchrotron radiation."467 lt ds however unclear why there should be such a correlation between the thermal far-infrared emission. and the non-thermal radio emission over such a wide range of galaxy types ancl masses. from starburst svstenis to more normal galaxies.," It is however unclear why there should be such a correlation between the thermal far-infrared emission and the non-thermal radio emission over such a wide range of galaxy types and masses, from starburst systems to more normal galaxies."468 As a result of this. many models resort toa relatively significant amount of fine tuning. such as assuming a much stronger magnetic field than what is estimated. via minimum energy arguments (e.g.Thompsonetal.2006).," As a result of this, many models resort to a relatively significant amount of fine tuning, such as assuming a much stronger magnetic field than what is estimated via minimum energy arguments \citep[e.g.][]{Thompson06}."469".. A full discussion of such arguments is bevond the scope of this paper: however. we refer to the reader to Lackietal.(2009) and Lacki&""Thompson(2009) who provide a detailed discussion of the various physical interpretations of the [ar-infrared.racio correlation (Εν)."," A full discussion of such arguments is beyond the scope of this paper; however, we refer to the reader to \citet{Lacki1} and \citet{Lacki2} who provide a detailed discussion of the various physical interpretations of the far-infrared--radio correlation (FIRC)."470 Observationally. recent work has concentrated on exploring the FIRC as a function of redshift. mainly. because of the preponderance of deep.Spifzer ancl radio data over relatively small <10 square degree areas (c.g.Appletonetetal. 2010).," Observationally, recent work has concentrated on exploring the FIRC as a function of redshift, mainly because of the preponderance of deep and radio data over relatively small $<10$ square degree areas \citep[e.g.][]{Appleton04, Frayer06, Ibar08, Murphy09,Michalowski10, Sargent10,Bourne10}."471. Phis has led to several authors suggesting hat 1e FIRC remains unchanged out to high redshift (2221.5) (c.g.Sargentctal.2010)... whereas others suggest a shallow ecrease in the ratio between far-infrared. luminosity and radio Luminosity (e.g.Sevmourctal.2009).," This has led to several authors suggesting that the FIRC remains unchanged out to high redshift $z\gtsim 1.5$ ) \citep[e.g.][]{Sargent10}, whereas others suggest a shallow decrease in the ratio between far-infrared luminosity and radio luminosity \citep[e.g.][]{Seymour09}."472.. Constraining rw evolution of the FIRC is important as it may clp oi understanding of the physical mechanism which results in such a tight correlation between the thermal and non- emission., Constraining the evolution of the FIRC is important as it may help our understanding of the physical mechanism which results in such a tight correlation between the thermal and non-thermal emission.473 For example. metallicity and temperature," For example, metallicity and temperature"474back the orbit of to show that a αναισα] disk runaway eveut ago is very likely the ejection mechanisi in this case.,back the orbit of to show that a dynamical disk runaway event ago is very likely the ejection mechanism in this case.475 has a total velocity referred to the local standard of vest (LSR) of (Aaitzenetal.1998) inakiug it the second fastest runaway star after the massive B-type giautIID271791., has a total velocity referred to the local standard of rest (LSR) of \citep{maaap339} making it the second fastest runaway star after the massive B-type giant.476. The various sinularitics between both stars were motivation to us to re-investieate the origin of60350., The various similarities between both stars were motivation to us to re-investigate the origin of.477. To this ai we carried out a quantitative analysis of a lugh-resolition spectra using non-local thermodyuanic equilibrimu (NLTE) techniques for the first time., To this aim we carried out a quantitative analysis of a high-resolution spectrum using non-local thermodynamic equilibrium (NLTE) techniques for the first time.478 Stellar paramctors were thus revised and elemental abundances constrained (Sect. 22))., Stellar parameters were thus revised and elemental abundances constrained (Sect. \ref{spectroscopy}) ).479 The results together with proper motions from the new reduction of the Catalog were used to determine the current three-dineusional (3D) space velocity., The results together with proper motions from the new reduction of the Catalog were used to determine the current three-dimensional (3D) space velocity.480 The following kincmatic analysis sugeested that the star originated iu or near the Crux-Scutiun spiral armi (Sect. ?7))., The following kinematic analysis suggested that the star originated in or near the Crux-Scutum spiral arm (Sect. \ref{kinematics}) ).481 Finally we discuss the kinematic paramucters and the clemenutal abuudauce pattern in the liebt of the rivaling formation scenarios. ilc. binary supernova versus dynamical cluster ejection (Sect. ?7)).," Finally we discuss the kinematic parameters and the elemental abundance pattern in the light of the rivaling formation scenarios, i.e., binary supernova versus dynamical cluster ejection (Sect. \ref{discussion}) )."482 was observed in 2008 December with he high-resolution echelle spectroeraph of the 9.2 wan Tobby-Eberly Telescope (ITET) at the McDonald Observatory., was observed in 2008 December with the high-resolution echelle spectrograph of the $9.2$ m Hobby-Eberly Telescope (HET) at the McDonald Observatory.483 Three individual spectra with resolving over A/AA=15000 and useful wavelength rangeAA.. were co-added. resulting in a signal-to-noise ratio AAJ](S/N) avounud 110 in the blue visual range.," Three individual spectra with resolving power $\lambda/\Delta \lambda=15\,000$ and useful wavelength range, ] were co-added, resulting in a signal-to-noise ratio (S/N) around $140$ in the blue visual range."484 Additionally. hree intermecdiate-resolution spectra taken in 2009 May and July with the 3.5 uuu telescope at Calar Alto. Spain. and its loneslt TWIN spectrograph were available chlareing the spectral coverage down toAA... naling accessible the high-order Baliner lines aud the Balmer jump.," Additionally, three intermediate-resolution spectra taken in 2009 May and July with the $3.5$ m telescope at Calar Alto, Spain, and its long-slit TWIN spectrograph were available enlarging the spectral coverage down to, making accessible the high-order Balmer lines and the Balmer jump."485 The quantitative spectral analysis was carried out following the hybrid NLTE approach discussed by Nieva&Przvbilla(2006.2007.2008) aud Przvbillaetal. (2006): liue-blanketed LTE 1inodel atinosplieres were computed with ATLAS9 (νο 1993).. while NLTE liue formation calculations were performiecl using updated versions of DETAIL and SURFACE (Cuddings1981:Butler&Giddings 1985).," The quantitative spectral analysis was carried out following the hybrid NLTE approach discussed by \citet{niapjl639, niaap467, niaap481} and \citet{praap445}: line-blanketed LTE model atmospheres were computed with ATLAS9 \citep{ku93}, , while NLTE line formation calculations were performed using updated versions of DETAIL and SURFACE \citep{giphd, bu9}. ."486. State-ofthe-art model atoms were adopted allowing absolute elemental abuudauces to be obtained with hieh accuracy (κος. Przvbillaetal.2008a.) for anu overview).," State-of-the-art model atoms were adopted allowing absolute elemental abundances to be obtained with high accuracy (see \citealt{prapjl688,prapj684} for an overview)."487 Atinospheric parameters and eleiieutal abundauces were derived by detailed line-profile analysis aud fitting of the spectral euergy distribution (SED)., Atmospheric parameters and elemental abundances were derived by detailed line-profile analysis and fitting of the spectral energy distribution (SED).488 The fundamental atmospheric paranieters0.15. incroturbuleut velocity and projected rotational velocity were primarily coustrained— from Baluer and lines as well as the ionization equilibrimm.," The fundamental atmospheric parameters, microturbulent velocity and projected rotational velocity were primarily constrained from Balmer and lines as well as the ionization equilibrium."489 Elemental abundances were then obtained by matching the nieasurable lines of the individual chemical species while keeping all other stellar parameters fixed (see Fig. 1))., Elemental abundances were then obtained by matching the measurable lines of the individual chemical species while keeping all other stellar parameters fixed (see Fig. \ref{abundance_uncertainties}) ).490 Tn the eud. a final svuthetic spectrum was computed which excellently reproduces the observation (sec Fie. 2)).," In the end, a final synthetic spectrum was computed which excellently reproduces the observation (see Fig. \ref{comparison}) ),"491 coufinüuug the B-type nature of60350., confirming the B-type nature of.492. Tuterestinely. a helm abundance higher than solar.=11.21.. was required to match the Ποια lines. the depth of the Daliuer lines aud the SED sinultaneouslv (see Fig. 3)).," Interestingly, a helium abundance higher than solar, was required to match the helium lines, the depth of the Balmer lines and the SED simultaneously (see Fig. \ref{sed}) )."493 The resulting abuudances (averaged over all lines of au clement) are τος iu Table 1.., The resulting abundances (averaged over all lines of an element) are listed in Table \ref{stellar}.494 spectra vielded a barvceutric radial velocityAll of|. equivalent to iu very eood agreement with deDoerctal.(1988) who fouud|.," All spectra yielded a barycentric radial velocity of, equivalent to in very good agreement with \citet{deaap202} who found."495 Bearing in mind the different time intervals between cach measurement. the star is unlikely a binary.," Bearing in mind the different time intervals between each measurement, the star is unlikely a binary."496 Comparing the location of in the (Tig.logg) diagram to evolutionary tracks (Schalleretal.1992) of solar inetallicity as shown in Fig.," Comparing the location of in the $\left(T_{\rm eff},\log g \right)$ diagram to evolutionary tracks \citep{scaaps96} of solar metallicity as shown in Fig."497 [allowed its mass and age to be constrained., \ref{evolution} allowed its mass and age to be constrained.498 The distance to could then be calculated from AL. V. Tog. logg aud extinction Ay=VV)=O0.07mae using the method described by Raispecketal.(2001) to bekpc.," The distance to could then be calculated from $M$, $V$, $T_{\rm eff}$, $\log g$ and extinction $A_V=3.1\,E(B-V)=0.07\,\rm mag$ using the method described by \citet{raaap379} to be."499. The precision of the analysis was restricted by an interplav of three effects: Ij Les in a telpcrative region where the opticalTWIP spectrum shows very few strone but nium weak metal lines., The precision of the analysis was restricted by an interplay of three effects: I) lies in a temperature region where the optical spectrum shows very few strong but many weak metal lines.500 II) A considerable fraction ofthelatter is s11icared out due to the high projected rotational velocity esiu/., II) A considerable fraction ofthelatter is smeared out due to the high projected rotational velocity .501. IIT) The, III) The5021997)]] to —2 [the singular isothermal sphere (SIS) case (Gott&Gunn1974:Gott1984)]] while maintaining the same mass density in lenses. the integral lensing probability increases by. more (han two orders of magnitudes lor the flat model of the Universe (LOO2).,"] to $-2$ [the singular isothermal sphere (SIS) case \citep{got74,tur84}] ] while maintaining the same mass density in lenses, the integral lensing probability increases by more than two orders of magnitudes for the flat model of the Universe (LO02)."503 Therefore. lensing also sensitivelv probes small scale structure.," Therefore, lensing also sensitively probes small scale structure."504 This complicates matters and renders it is hazardous (ο use observed lensing statistics to draw inferences with regard to cosmology before determining the sensitivity to other factors., This complicates matters and renders it is hazardous to use observed lensing statistics to draw inferences with regard to cosmology before determining the sensitivity to other factors.505 In LOQ2. we have shown that in order to explain the observed numbers of lenses found in the JVAS/CLASS survey. at least. (wo populations of dark halos must exist in nature.," In LO02, we have shown that in order to explain the observed numbers of lenses found in the JVAS/CLASS survey, at least two populations of dark halos must exist in nature."506 CNme population. which corresponds to normal galaxies. has masses <10M. and a steep inner densitv profile (azz2. i.e. SIS) presumably determined by the distribution of baryonic natevial in the inner parts of the other one. which corresponds (0 groups or clusters of galaxies. has masses >10!M. and a shallow inner density profile (a1.4. ie. similar to NEW).," One population, which corresponds to normal galaxies, has masses $\la 10^{13} M_\odot$ and a steep inner density profile $\alpha \approx 2$, i.e. SIS) presumably determined by the distribution of baryonic material in the inner parts of the other one, which corresponds to groups or clusters of galaxies, has masses $\ga 10^{13} M_\odot$ and a shallow inner density profile $\alpha \la 1.4$, i.e. similar to NFW)."507 A similar conclusion has been obtained by Porciani&Macau(2000). [or explaining the number of lenses found in the CASTLES survey., A similar conclusion has been obtained by \citet{por00} for explaining the number of lenses found in the CASTLES survey.508 These results are consistent with the theoretical studies on the cooling of massive gas clouds: there is a critical mass of halos ~10M. below which cooling of the corresponding barvonic component will lead to concentration of the barvous to the inner parts of the mass profile al. 1986).," These results are consistent with the theoretical studies on the cooling of massive gas clouds: there is a critical mass of halos $\sim 10^{13} M_\odot$ below which cooling of the corresponding baryonic component will lead to concentration of the baryons to the inner parts of the mass profile \citep{ree77,blu86}."509. In this paper we investigate (he lensing statistics produced by a compound population of halos., In this paper we investigate the lensing statistics produced by a compound population of halos.510" We assume (hat there are three populations of halos in (he Universe: Population A: LOMATAL.<Mc:[E]X--—-=crusmes .à—2 (SIS): Population B: M>2x10""hHAL. a=1.3 [GNFW (generalizedNEW.Zhao 1996)]]: Population C: AZ<10hTAL... a=L3 (GNEW). where { is the IIubble constant in units of 100 kms ! |."," We assume that there are three populations of halos in the Universe: — Population A: $10^{10} h^{-1} M_\odot < M < 2\times 10^{13} h^{-1} 511M_\odot$ , $\alpha = 2$ (SIS); — Population B: $M > 2\times 10^{13} h^{-1} M_\odot$, $\alpha = 1.3$ [GNFW \citep[generalized NFW,][]{zha96}] ]; — Population C: $M < 10^{10} h^{-1} M_\odot$, $\alpha = 1.3$ (GNFW), where $h$ is the Hubble constant in units of 100 km $^{-1}$ $^{-1}$."512 Population A corresponds Lo spiral ancl elliptical galaxies. whose centers are dominated by baryonic matter.," Population A corresponds to spiral and elliptical galaxies, whose centers are dominated by baryonic matter."513 Population D corresponds to groups or clusters of galaxies. whose centers are dominated by dark matter.," Population B corresponds to groups or clusters of galaxies, whose centers are dominated by dark matter."514 Population C corresponds todwarl galaxies or subgalactie objects. whose centers lack barvons," Population C corresponds todwarf galaxies or subgalactic objects, whose centers lack baryons"515When we increase the 1tunber of the Gauss-Seidel iteration during a cycle of the inulti-grid iteration. we get a higher recuctiou in the residual per cycle at the expeuse of lounger computation ime.,"When we increase the number of the Gauss-Seidel iteration during a cycle of the multi-grid iteration, we get a higher reduction in the residual per cycle at the expense of longer computation time."516 The recluction per uuii| colnputaion load is higher when we the Catss-Seicdel iteration is »erforijecd several times each., The reduction per unit computation load is higher when we the Gauss-Seidel iteration is performed several times each.517 We tried the successive over relaxaion (SOR) itsteacl the Causs-Seicdel iteration to accelerate he converge but failec., We tried the successive over relaxation (SOR) instead the Gauss-Seidel iteration to accelerate the converge but failed.518 Wren SOR ts used as the pre- aud post-relaxation in our imulti-ericl iteration. the residual iucreases sometiues cdepeucii[n]0 ou the initial guess.," When SOR is used as the pre- and post-relaxation in our multi-grid iteration, the residual increases sometimes depending on the initial guess."519 SOR works well ouly when use one level of the grid. Le.. when the grid is iol nestec.," SOR works well only when use one level of the grid, i.e., when the grid is not nested."520 We also tried to improve the interpolatiou formia. Equation (19)). lor a higher order accuracy.," We also tried to improve the interpolation formula, Equation \ref{interpolation1}) ), for a higher order accuracy."521 We [οιud that our iter:ion did uot converge whet we used a higher order interpolation formula., We found that our iteration did not converge when we used a higher order interpolation formula.522 Thougl we Ceui not εἶοιy existence of a successful i1erpolation formula. we could uot find it.," Though we can not deny existence of a successful interpolation formula, we could not find it."523 We evaluated the computation oad of our multi-erid algorithin by measuring the computation time., We evaluated the computation load of our multi-grid algorithm by measuring the computation time.524 A UNIX workstation. SGI 05 (MIPS R10000 250 MHz) was used for the measurement.," A UNIX workstation, SGI O2 (MIPS R10000 250 MHz) was used for the measurement."525 The computation iine was Lnueasured with the subroutine DTIME., The computation time was measured with the subroutine DTIME.526 Figure 9 shows the execution imme per FMG evcle., Figure \ref{etime.eps} shows the execution time per FMG cycle.527" When Tje abscissa deuotes the effective total number. Noell = finas""+ dnde logarithinie scale."," When The abscissa denotes the effective total number, $ N _{\rm cell} $ = $528\ell _{\rm max} \, N ^3 $, in the logarithmic scale."529 The ordinate denotes the logarithia of the computation liive. /.," The ordinate denotes the logarithm of the computation time, $ t $."530 The dashed lies clenote the relajons. CPU ime xοι and CPU time.," The dashed lines denote the relations, CPU time $ \propto \, N _{\rm cell} $ and CPU time."531" We uade this pkX |wy chaugiug (gas al aN in the range of 2""dí10 and 5«ONx«12k."," We made this plot by changing $ \ell _{\rm max} $ and $ N $ in the range of $ 2 \, \le \, \ell \, \le \, 10 $ and $ 8 \, \le \, N532\, \le \, 128 $."533 As slOW Lin Figure 9.. the coriputation load is p'oportional o the Αι when Noο32.," As shown in Figure \ref{etime.eps}, the computation load is proportional to the $ N _{\rm cell} $ when $ N \, \ge \, 32 $."534 This Ineaus thi ilour aleorithin is scalabe for ποται auc arge nested erids., This means that our algorithm is scalable for medium and large nested grids.535" WeaSO lleastl'ed the computation time with the super compuer Fujitsu VPP5000 at National Astronomica Obse""alory. Japau."," We also measured the computation time with the super computer Fujitsu VPP5000 at National Astronomical Observatory, Japan."536" When N=256 a1| {Wax5. the CPU time is 0.21 sec per FMCG evele,"," When $ N \, = \, 256 $ and $ \ell _{\rm max} \, = \, 5 $, the CPU time is 0.21 sec per FMG cycle."537 This CPU time is reasonably small σοιiparecl with the CPU time for solving the hivcdrodsnamical eqlation on the same nested grid., This CPU time is reasonably small compared with the CPU time for solving the hydrodynamical equation on the same nested grid.538 Tlis CPU time is uot scalable to the number of cells siuce the suyer computer has vector processors aud its CPU time is not proportional to the computation load., This CPU time is not scalable to the number of cells since the super computer has vector processors and its CPU time is not proportional to the computation load.539 As shown in the previous section. our numerical method provides an accurate solution with a reasonably small computation cost.," As shown in the previous section, our numerical method provides an accurate solution with a reasonably small computation cost."540 The computation load is scalable in the seuse that it is proportional to the uumber of the cells contained iu the nested grid., The computation load is scalable in the sense that it is proportional to the number of the cells contained in the nested grid.541 Our discrete Poissou equatiou is robust iu the sense that it cau be applied also to AMI as far as a parent cell is subcdivided into, Our discrete Poisson equation is robust in the sense that it can be applied also to AMR as far as a parent cell is subdivided into542with the highest overdensity particle. we strouud each potential deusitv maxima by a sphere of radius ring=105.!spe aud exclude all particles within this sphere from further search.,"with the highest overdensity particle, we surround each potential density maximum by a sphere of radius $r_{\rm find}=10h^{-1}\ \rm kpc$ and exclude all particles within this sphere from further search."543 The search radius is defined by the size of sunallest svstenis we aiu to ileutifv., The search radius is defined by the size of smallest systems we aim to identify.544 We verified that the results do not chanee if this radius is decreased by a factor ofup to four., We verified that the results do not change if this radius is decreased by a factor of up to four.545 After all potential halo centers ave identified. we the deusity distribution aud velocitiesof surrounding particles analyzeto test whethercenter corresponds to a bound chup.," After all potential halo centers are identified, we analyze the density distribution and velocities of surrounding particles to test whether the center corresponds to a gravitationally bound clump."546 Specifically. thewe coustruct the density. circular velocity. and velocity dispersion profiles around cach ceuter and iteratively remove wnbotud particlesdetails).," Specifically, we construct the density, circular velocity, and velocity dispersion profiles around each center and iteratively remove unbound particles."547 We then coustruct final profiles using ouly bound particles and use them to calculate such halo properties as the naxinuun circular velocity Vj. mass M. ete.," We then construct final profiles using only bound particles and use them to calculate such halo properties as the maximum circular velocity $V_{\rm m}$, mass $M$, etc."548 The virial radius is 11020inegless for the sublialos within a arecr host as their outer lavers are tidally stripped aud the extent of the halo is truncated., The virial radius is meaningless for the subhalos within a larger host as their outer layers are tidally stripped and the extent of the halo is truncated.549 The definition of the outer voundary of a subhalo and its mass are thus somewhat aüubieuous., The definition of the outer boundary of a subhalo and its mass are thus somewhat ambiguous.550 We adopt theradius. ἐν. at which he logarithmic slope of the deusitv profile constructed roni the bound particles becomes larger than 0.5 as we do not expect the density profile of the CDM. halos o be flatter than this slope.," We adopt the, $r_{\rm t}$, at which the logarithmic slope of the density profile constructed from the bound particles becomes larger than $-0.5$ as we do not expect the density profile of the CDM halos to be flatter than this slope."551 Empirically. this definition roughly correspouds to the radius at which the density of he eravitationally bound particles is equal to the backeround rost halo density. albeit witha large scatter.," Empirically, this definition roughly corresponds to the radius at which the density of the gravitationally bound particles is equal to the background host halo density, albeit witha large scatter."552 For some walos rg is larger than their virial radius., For some halos $r_{\rm t}$ is larger than their virial radius.553" Iu this case. we seta,—Rey."," In this case, we set $r_{\rm t}=R_{\rm vir}$."554 Throughout this paper. we will denote the ΠΕ of the virial mass and mass within r. simply as AM.," Throughout this paper, we will denote the minimum of the virial mass and mass within $r_{\rm t}$, simply as $M$."555 For ΠΕeach halo we also construct the circular velocity xofile V.tr)=ναλέςΕν aud compute the maximi circular velocity profile V4., For each halo we also construct the circular velocity profile $V_{c}(r)=\sqrt{GM(<r)/r}$ and compute the maximum circular velocity profile $V_{\rm m}$.556 Figue 2. shows the particle distribution iu the halo oat 2=0 along with the halos (circles) ideutified bv the halo finder., Figure \ref{fig:dmh} shows the particle distribution in the halo $_1$ at $z=0$ along with the halos (circles) identified by the halo finder.557 The particles are color-coded on a erav scale according to the logarithim of their density to enliance visibility of substructure clumps., The particles are color-coded on a gray scale according to the logarithm of their density to enhance visibility of substructure clumps.558" The radius of the largest πο indicates the actual virial radius. a. of the host gravitationally (Ry,=208h! kpe) the radii of the other halos are the münuuuni of the truncation radius rj aud ZA."," The radius of the largest circle indicates the actual virial radius, $R_{\rm vir}$, of the host halo $R_{\rm vir}=298h^{-1}$ kpc); the radii of the other halos are the minimum of the truncation radius $r_{\rm t}$ and $R_{\rm vir}$."559 The figure demonstrates that the algorithin is efficient iu identifving the substructure down to suiidll masses., The figure demonstrates that the algorithm is efficient in identifying the substructure down to small masses.560 The halo finder described above was run at the 96 saved epochs between 2=10 aud 2=0 with a typical spacing of 1]2.4105 vr between outputs., The halo finder described above was run at the 96 saved epochs between $z=10$ and $z=0$ with a typical spacing of $\sim 1-2\times 10^8$ yr between outputs.561" For cach epoch. the halo finder produced a halo catalog with positions. velocities. radi my,=dare.eee). nasses 20rng). maxima of the circular velocity profile V4, aud the radius at which the WANT occurs Maas."," For each epoch, the halo finder produced a halo catalog with positions, velocities, radii $r_{\rm h}=\min (r_{\rm562 t},r_{\rm vir})$, masses $m(<r_{\rm h})$, maximum of the circular velocity profile $V_{\rm m}$ and the radius at which the maximum occurs $r_{\rm max}$."563 du addition. for each halo we save indices of all eravitationally-bound DAL particles located within ng.," In addition, for each halo we save indices of all gravitationally-bound DM particles located within $r_{\rm h}$."564 This information is used to ideutifv the progenitors of halos at successive epochs., This information is used to identify the progenitors of halos at successive epochs.565 Specifically. for a current epoch u.s starting at +=0. we search progenitors for cach halo at several previous epochs z; as follows.," Specifically, for a current epoch $z_{\rm i}$, starting at $z=0$, we search progenitors for each halo at several previous epochs $z_{\rm i-j}$ as follows."566 First. we select a given fraction. found. of the most bound particles of the halos at the epochs of cousideration.," First, we select a given fraction, $f_{\rm bound}$, of the most bound particles of the halos at the epochs of consideration."567 We then compare the fraction of these particles that is common between all pairs of halos at successive epochs and assume that the halo with the highest common fraction is the progenitor., We then compare the fraction of these particles that is common between all pairs of halos at successive epochs and assume that the halo with the highest common fraction is the progenitor.568 The trajectories used in this study were constructed using froma= 0.25., The trajectories used in this study were constructed using $f_{\rm bound}=0.25$ .569 As the halo catalogs iav. miss some halos. especially near the completeness limit of the simulation. if the progenitor is not found at the previous epoch we," As the halo catalogs may miss some halos, especially near the completeness limit of the simulation, if the progenitor is not found at the previous epoch we"570where Bis the magnetic field at the pole. R is the radius of the star. ο) is the angular rotational frequeucyv of the star. e is the speed of light. » is the magnetic braking iudex.,"where $B$ is the magnetic field at the pole, $R$ is the radius of the star, $\Omega$ is the angular rotational frequency of the star, $c$ is the speed of light, $n$ is the magnetic braking index."571 For an aligued rotator without field decay. the braking index is roughly ο~3. however due to magnetic flux expulsion from a CFL QS. the magnetic field decays as prescribed by Nicbergaletal.(2006).," For an aligned rotator without field decay, the braking index is roughly $n\sim 3$, however due to magnetic flux expulsion from a CFL QS, the magnetic field decays as prescribed by \cite{niebergal06}."572. This results in an evolution of the lunineosity due to spiu-down. which is expressed by the relation. where the characteristic spiu-down tine (iu seconds) ls. Tu the above equatious. Moa is the QS mass. £4 is the initial spin period. and By is the initial maguetic field streneth.," This results in an evolution of the luminosity due to spin-down, which is expressed by the relation, where the characteristic spin-down time (in seconds) is, In the above equations, $M_{\rm QS}$ is the QS mass, $P_{0}$ is the initial spin period, and $B_{0}$ is the initial magnetic field strength."573 From Eq., From Eq.574 2 0ie can see that the hnuuinosity. due to rotational cucrev extracted from spin-down of a QS. has a natural break at time 7.," \ref{eq:qssd_lum} one can see that the luminosity, due to rotational energy extracted from spin-down of a QS, has a natural break at time $\tau$."575 Thus. if there was a oue to one relationship between spin-down Iuuinosity aud observed enission. then the power law decay of the observed lisht-curve should change from zero to 5/3 after roughly ten thousaud seconds.," Thus, if there was a one to one relationship between spin-down luminosity and observed emission, then the power law decay of the observed light-curve should change from zero to $-5/3$ after roughly ten thousand seconds."576 However. the observed ciission nuelit be modified by the forward shock as discussed iu Panaiteseu(2007).," However, the observed emission might be modified by the forward shock as discussed in \citet{panaitescu07}."577 The cnerey released from the sou-down of the QS is likely to be in the form of an ele wind., The energy released from the spin-down of the QS is likely to be in the form of an $e^+e^-$ wind.578 Thus. it should be mostly barvon free. since the QS becomes bare innuecdiately following its birth as it enters the CFL phase (seeNicherealetal.2006).," Thus, it should be mostly baryon free, since the QS becomes bare immediately following its birth as it enters the CFL phase \citep[see][]{niebergal06}."579. As in the case of a pulsar. spin-down euergv extracted from a QS is mainly in the equatorial plane.," As in the case of a pulsar, spin-down energy extracted from a QS is mainly in the equatorial plane."580 Bucciautinietal.(2007) performed miuerical παπαΊος where they showed that it is still vossible to collimate such equatorial flows iuto a jet., \citet{bucciantini07} performed numerical simulations where they showed that it is still possible to collimate such equatorial flows into a jet.581 A relativistic outflow from the spin-down of a hielilv-naguetized star has been suggested before as a uechanisui to produce plateaus (forinstanceinTrojaetal. 2007).. however they did not propose a unified model explaining both the prompt enussiou aud the afterglow catures.," A relativistic outflow from the spin-down of a highly-magnetized star has been suggested before as a mechanism to produce plateaus \citep[for instance in][]{troja07}, however they did not propose a unified model explaining both the prompt emission and the afterglow features."582 We have here proposed a model that cau explain oth the prompt CRB cunission aud the observed N-rav afterelow features., We have here proposed a model that can explain both the prompt GRB emission and the observed X-ray afterglow features.583 Eq., Eq.584 2 naturally gives a break in the cneime Iuninositv at fτ., \ref{eq:qssd_lum} naturally gives a break in the engine luminosity at $t=\tau$.585 The cneine will also remain active after this weak. but the eugine huninosity will eraduallv dec: (with a power law ~—5/3:56k which is not necessarily he power law decay in the observed emission).," The engine will also remain active after this break, but the engine luminosity will gradually decay (with a power law $\sim -5/3$; which is not necessarily the power law decay in the observed emission)."586 Iu some instances however. it is possible that the QS reaches au uustable configuration. such that the QS stage is oulv eniporary before the collapse to a DII.," In some instances however, it is possible that the QS reaches an unstable configuration, such that the QS stage is only temporary before the collapse to a BH."587 If the OS collapses to a BI during spin-down. the eneiue will likely be shut off.," If the QS collapses to a BH during spin-down, the engine will likely be shut off."588 Although the BIT is likely to )o rapidly rotating. a disk 1s necessary in order to extract he rotational cucrey of a BIT through the Dlaudford-Zuajek iechanisi (BZ:Dlaudford&Zuajel1977).," Although the BH is likely to be rapidly rotating, a disk is necessary in order to extract the rotational energy of a BH through the Blandford-Znajek mechanism \citep[BZ;][]{bz77}."589 Ouly if a disk has remained around the QS during down or if it is formed after the formation of the DIT. cau the BZ inechanisin play a vole.," Only if a disk has remained around the QS during spin-down or if it is formed after the formation of the BH, can the BZ mechanism play a role."590 If this does not occur. the observed light curve will be generated by the external shock only after this stage.," If this does not occur, the observed light curve will be generated by the external shock only after this stage."591 A sharp drop off will be secu as the light curve drops from the level given by the spin-down outflow to the level given by the external shock., A sharp drop off will be seen as the light curve drops from the level given by the spin-down outflow to the level given by the external shock.592 We sugeestOO that iu CRB μον curves exhibiting plateaus. those possessing a eradual decay following the plateau are either duc to rereshed shocks as discussed iu SOBO?7 or from spin-down of OSs that have not collapsed to BI," We suggest that in GRB light curves exhibiting plateaus, those possessing a gradual decay following the plateau are either due to refreshed shocks as discussed in SOB07 or from spin-down of QSs that have not collapsed to BHs."593 If the secondary ouflow is responsible for the X-rav afterelow.s. then the external shock cau produce the optical afterelow.," If the secondary outflow is responsible for the X-ray afterglow, then the external shock can produce the optical afterglow."594 This sceiuio nüeht explain why the optical aud X-rav afterelows behave different iu some CRBs., This scenario might explain why the optical and X-ray afterglows behave different in some GRBs.595 Iu this section we will apply our model to two CRBs. GRDB 070110 and CRB oGOG07TA. that both show a Hattening followed by a sharp drop off which is difficult o explain with the external shock.," In this section we will apply our model to two GRBs, GRB 070110 and GRB 060607A, that both show a flattening followed by a sharp drop off which is difficult to explain with the external shock."596 Some observed xoperties of both GRBs are sununarized in Table 1.., Some observed properties of both GRBs are summarized in Table \ref{obstable}.597 Based ou observations of the duration of the N-ray Hattenimg. we use Eq.," Based on observations of the duration of the X-ray flattening, we use Eq."598 3) to estimate the correspoudine uagnetic field streneth, \ref{eq:charac_time} to estimate the corresponding magnetic field strength.599 We then use Eq., We then use Eq.600 2. to find he spin-down huuinosity., \ref{eq:qssd_lum} to find the spin-down luminosity.601 Both the maguetic field aud he spin-down lwiuinosity found this wav are listed iu Table 2. which is then compared to observed values (Table 13)., Both the magnetic field and the spin-down luminosity found this way are listed in Table \ref{calculatedtable} which is then compared to observed values (Table \ref{obstable}) ).602 Furthermore. now that we have au estimate or the mmaenetic field of the QS. this eives us an estimate or the accretion rate that can be channeled to the polar cap.," Furthermore, now that we have an estimate for the magnetic field of the QS, this gives us an estimate for the accretion rate that can be channeled to the polar cap."603 We asstune a jet opening augle of about 10 degrees., We assume a jet opening angle of about 10 degrees.604 The observed prompt CRB ciuission is then calculated * assundüues that a combination of accretion cficicucy and radiative cfiiciency leads to ~| of the total eravitational energy of the accreted material is converted o prompt radiation., The observed prompt GRB emission is then calculated by assuming that a combination of accretion efficiency and radiative efficiency leads to $\sim 1\%$ of the total gravitational energy of the accreted material is converted to prompt radiation.605 As shown below. for both. CRB VFOLLO aud CRB 060607À we find that the maguectic ficld found based ou the duration of the N-vav flattcnine consistently and simmitancously explains the euerev of voth the GRB itself aud the Nay flattening.," As shown below, for both GRB 070110 and GRB 060607A we find that the magnetic field found based on the duration of the X-ray flattening consistently and simultaneously explains the energy of both the GRB itself and the X-ray flattening."606with an identical equation for dco/dt.,with an identical equation for $d\varpi_k/dt$.607" Note that in this formulation, as before, the secular term is linearly proportional to sin(/), unlike the GR. correction."," Note that in this formulation, as before, the secular term is linearly proportional to $\sin(I)$, unlike the GR correction."608" Thus, the system inclination can be solved for in the same way as above, but without constraints on eccentricity of semi-major axes."," Thus, the system inclination can be solved for in the same way as above, but without constraints on eccentricity of semi-major axes."609" 'To date, the number of detected multi-planet systems that host small close-in planets remains limited to a handful of systems: HD 40307, 55 Cnc, 61 Vir, GJ 581 and GJ 876."," To date, the number of detected multi-planet systems that host small close-in planets remains limited to a handful of systems: HD 40307, 55 Cnc, 61 Vir, GJ 581 and GJ 876."610" Furthermore, the data for these systems are still comparatively sparse, so the error bars on the planet’s eccentricities are rather large."," Furthermore, the data for these systems are still comparatively sparse, so the error bars on the planet's eccentricities are rather large."611" These issues will surely get resolved with time, but at this point we can only give a rough assessment, and shall limit our analysis to a single case: 61 Vir."," These issues will surely get resolved with time, but at this point we can only give a rough assessment, and shall limit our analysis to a single case: 61 Vir."612 'The planetary system around the nearby sun-like star 61 Vir was discovered by Vogt et al (2010)., The planetary system around the nearby sun-like star 61 Vir was discovered by Vogt et al (2010).613" The star hosts 3 planets, with orbital periods of roughly 4.2d, 38d and 124d (see Table 1 for an orbital fit)."," The star hosts 3 planets, with orbital periods of roughly $4.2$ d, $38$ d and $124$ d (see Table 1 for an orbital fit)."614 A simple evaluation of the system's dynamical stability yields no useful constraints on the inclination of the system., A simple evaluation of the system's dynamical stability yields no useful constraints on the inclination of the system.615" However, the minimum mass of the inner-most planet of m=5.1E corresponds to that of a super-Earth, making it an 0.6Mgideal candidate for our method."," However, the minimum mass of the inner-most planet of $\tilde{m}=5.1 \pm 0.6 M_{\oplus}$ corresponds to that of a super-Earth, making it an ideal candidate for our method."616 'The characteristic isolated circularization timescale of planet b is roughly τεQx10° years., The characteristic isolated circularization timescale of planet b is roughly $\tau_c \sim Q \times 10^6$ years.617" A damped, modified LL solution in Figure 3) reveals that depending on starting (shownconditions, up to 10 7 is required for the system to arrive to the fixed point."," A damped, modified LL solution (shown in Figure 3) reveals that depending on starting conditions, up to 10 $\tau$ is required for the system to arrive to the fixed point."618" Thus, as already pointed out by Vogt et al given the star's multi-billion year age, we expect the (2010),system to be stationary if Qy<10°. For the illustrative purposes of this paper, we assume that planet b's tidal quality factor is similar to that of rocky bodies i.e. Q,=100."," Thus, as already pointed out by Vogt et al (2010), given the star's multi-billion year age, we expect the system to be stationary if $Q_b \lesssim 10^3.$ For the illustrative purposes of this paper, we assume that planet b's tidal quality factor is similar to that of rocky bodies i.e. $Q_b = 100$."619" Initially, we proceed as described in section 2.1 and compute the surviving LL eigenvector that physically corresponds to a state where all orbits are apsidally aligned."," Initially, we proceed as described in section 2.1 and compute the surviving LL eigenvector that physically corresponds to a state where all orbits are apsidally aligned."620" Given the moderate eccentricity (e> of the outer two planets, however, the LL solution does 0.1)not give a quantitatively acceptable answer."," Given the moderate eccentricity $e > 0.1$ ) of the outer two planets, however, the LL solution does not give a quantitatively acceptable answer."621" Consequently, we recompute the eccentricity ratios using the Gaussian averaging method, as described in section 2.4, utilizing the LL solution as an initial guess in the root-finding algorithm."," Consequently, we recompute the eccentricity ratios using the Gaussian averaging method, as described in section 2.4, utilizing the LL solution as an initial guess in the root-finding algorithm."622 The resulting curves are plotted in figures (4) and (5)., The resulting curves are plotted in figures (4) and (5).623" It is noteworthy that although the Gaussian and LL solutions are qualitatively similar, higher-order secular terms clearly make a noticeable contribution to the fixed-point solution."," It is noteworthy that although the Gaussian and LL solutions are qualitatively similar, higher-order secular terms clearly make a noticeable contribution to the fixed-point solution."624" Although the error bars on the orbital elements are still large, it is noteworthy that the observed system is consistent with a fixed point configuration."," Although the error bars on the orbital elements are still large, it is noteworthy that the observed system is consistent with a fixed point configuration."625" Thus, further observation of the system is warranted, given that if the system is found to be in a stationary state, it would yield not only the true masses, but also a constraint on the tidal quality factor of the inner-most planet."," Thus, further observation of the system is warranted, given that if the system is found to be in a stationary state, it would yield not only the true masses, but also a constraint on the tidal quality factor of the inner-most planet."626" 'The domain of applicability of the method described in this paper does not extend to ""large"" planets that we require AP...)«1 in order to solve for (recall sin(I))"," The domain of applicability of the method described in this paper does not extend to “large"" planets (recall that we require $\Lambda_{tidal}^{p} \ll 1$ in order to solve for $\sin(I)$ )."627" However, for massive, close-in planets, the sin(I) degeneracy can be resolved from spectral characterization of the host star alone (Snellen et al 2010)."," However, for massive, close-in planets, the $\sin(I)$ degeneracy can be resolved from spectral characterization of the host star alone (Snellen et al 2010)."628" In such case, the orbital precession rate yields information on athe radius and the interior structure of the planet."," In such a case, the orbital precession rate yields information on the radius and the interior structure of the planet."629" If only a single planet is present in the system, then the method described by Ragozzine Wolf (2009) can be employed."," If only a single planet is present in the system, then the method described by Ragozzine Wolf (2009) can be employed."630" Namely, if the planet is sufficiently close to its host star, the orbital precession rate may be as high as a few "," Namely, if the planet is sufficiently close to its host star, the orbital precession rate may be as high as a few degrees/year."631"In this case, direct observation of the orbital precession degrees/year.can be related to the sum of equations (7) -(9)."," In this case, direct observation of the orbital precession can be related to the sum of equations (7) -(9)."632" As already discussed above, however, the first"," As already discussed above, however, the first"633of Be/X-rav binary pulsars are usually hard.,of Be/X-ray binary pulsars are usually hard.634 A fluorescent iron emission line at GA keV is observed in the spectrum of most of the X-ray pulsars., A fluorescent iron emission line at 6.4 keV is observed in the spectrum of most of the X-ray pulsars.635 Lt is possible that most. of these systems have a soft X-ray excess above the power-law continuum component., It is possible that most of these systems have a soft X-ray excess above the power-law continuum component.636 However. detection of the the soft excess depends on the value of absorption column density (Paul et al.," However, detection of the the soft excess depends on the value of absorption column density (Paul et al."637 2002: Naik Paul 2004a. 2004b and references therein).," 2002; Naik Paul 2004a, 2004b and references therein)."638 The transient N-rav pulsar was discovered. on 1993 July 1 by the BATSE experiment onboard the(CORO) (Stollberg ct al., The transient X-ray pulsar was discovered on 1993 July 14 by the BATSE experiment onboard the (Stollberg et al.639 1993)., 1993).640 X-ray pulsations of 93.587 s were detected in the 20-120 keV οποίον range of BAVSE., X-ray pulsations of 93.587 s were detected in the 20-120 keV energy range of BATSE.641 From ASCA observation. the X-ray. pulse profile of the pulsar was found to have a clouble-peak structure with a well-defined. sharp intensity minimum and a less prominent seconcary minimum (Tanaka et al.," From ASCA observation, the X-ray pulse profile of the pulsar was found to have a double-peak structure with a well-defined, sharp intensity minimum and a less prominent secondary minimum (Tanaka et al."642 19921.Το BATSE spectrum was described. by an optically thin thermal bremsstrahlung model withAY = 25 keV. Following the discovery. the optical ancl infrared observations of the optical counterpar o revealed the presence of strong Balmer emission lines and infrared excess (Coe et al.," 1993).The BATSE spectrum was described by an optically thin thermal bremsstrahlung model with = 25 keV. Following the discovery, the optical and infrared observations of the optical counterpart to revealed the presence of strong Balmer emission lines and infrared excess (Coe et al."643 1994)., 1994).644 Base on these results. the svstem was classified as a massive NALLY systeni consisting of a neutron star as the compac object and a Be or a supergiant primary.," Based on these results, the system was classified as a massive binary system consisting of a neutron star as the compact object and a Be or a supergiant primary."645 After 260 days of his outburst. a second outburst was detected by BATSE (Finger ct al.," After 260 days of this outburst, a second outburst was detected by BATSE (Finger et al."646 19904)., 1994).647 Assuming this 260 d as the orbita »eriod. of the pulsar. Finger et al. (," Assuming this 260 d as the orbital period of the pulsar, Finger et al. ("6481994). estimated. the mass of the binary companion to be 3-8M. indicating the system as a high mass X-rav binary.,1994) estimated the mass of the binary companion to be 3-8$M_\odot$ indicating the system as a high mass X-ray binary.649 The ROSAT PSPC observation of the pulsar. in the declining. phase of the discovery outburst by BATSE in 1993. clearly. detected the 93.4 s pulsation with a couble-peaked pulse profile in O.1-2.4 keV lieht curve (Petre Gehrels L994).," The ROSAT PSPC observation of the pulsar, in the declining phase of the discovery outburst by BATSE in 1993, clearly detected the 93.4 s pulsation with a double-peaked pulse profile in 0.1-2.4 keV light curve (Petre Gehrels 1994)."650 A search. in the archive of ENOSATYMedium-Energy Experiment (ALLE) observation. centered on LID SS661. revealed. the presence of the pulsed. emission. at. the same period curing the 1993 outburst (Macomb. Shrader. Schultz 1994).," A search in the archive of EXOSAT/Medium-Energy Experiment (ME) observation, centered on HD 88661, revealed the presence of the pulsed emission at the same period during the 1993 outburst (Macomb, Shrader, Schultz 1994)."651 The X-rav spectrum (0.510 keV range) was found to be highly absorbed (Ng=0.7LOalomsem 7) and. described bv a hard power-law with a photon index of ~1.2.," The X-ray spectrum (0.8–10 keV range) was found to be highly absorbed $N_H652= 0.7\times10^{22} atoms cm^{-2}$ ) and described by a hard power-law with a photon index of $\sim$ 1.2."653 A combined analysis of data from the CGRO and ASCA observations. though non-simultaneous. shortly. after. the peak of the discovery outburst. reported that the broadband spectrum of the pulsar can be well approximated by a power-law with an exponential cutolf and a 6.4 keV iron emission line (Shrader ct al.," A combined analysis of data from the CGRO and ASCA observations, though non-simultaneous, shortly after the peak of the discovery outburst, reported that the broadband spectrum of the pulsar can be well approximated by a power-law with an exponential cutoff and a 6.4 keV iron emission line (Shrader et al."654 1999)., 1999).655 The pulse profile was also found to be energy dependent. a couble-peaked profile detected by ASCA that evolved into a single-peaked. profile as detected by BATSE.," The pulse profile was also found to be energy dependent, a double-peaked profile detected by ASCA that evolved into a single-peaked profile as detected by BATSE."656 Analyzing the BATSE and Rossiraydiminglvcplorer(PNT E)/SM Εαν histories. Shracer et al. (," Analyzing the BATSE and $Rossi 657X-ray Timing Explorer (RXTE)$ /ASM flux histories, Shrader et al. ("6581999) suggested the orbital period of the svstem to be 7135 days.,1999) suggested the orbital period of the system to be $\sim$ 135 days.659 However. Levine Corbet (2006) detected a 248.9 day periodicity in the [NXTE Xll-sky Monitor (ASAI) X-ray light curve by analyzing data accumulated over nearly 10 vears.," However, Levine Corbet (2006) detected a 248.9 day periodicity in the $RXTE$ All-sky Monitor (ASM) X-ray light curve by analyzing data accumulated over nearly 10 years."660 This periodicity was found by visual identification of periodically occurring outbursts in the ASM light curve., This periodicity was found by visual identification of periodically occurring outbursts in the ASM light curve.661 An independent. analysis of pulse period. variations during outbursts. using DATSIS data. estimated the orbital period precisely. to be 247.8 d (Coe et al.," An independent analysis of pulse period variations during outbursts, using BATSE data, estimated the orbital period precisely to be 247.8 d (Coe et al."662 2007) which is good agreement with the orbital period. determined. from. the recurrence of the X-ray outbursts in ASAT light curve., 2007) which is good agreement with the orbital period determined from the recurrence of the X-ray outbursts in ASM light curve.663 Following the detection of an intense outburst [rom with the Burst Alert Telescope (BAT) on Swift on 2007 November 17 (Ixrimm ct 22007). the accreting N-rav pulsar was observed. with various. N-rav observatories.," Following the detection of an intense outburst from with the Burst Alert Telescope (BAT) on Swift on 2007 November 17 (Krimm et 2007), the accreting X-ray pulsar was observed with various X-ray observatories."664 The RATE observations detected the pulsar up to ~70 keV along with the regular 793.75 s pulsations (Wilms et al., The RXTE observations detected the pulsar up to $\sim$ 70 keV along with the regular $\sim$ 93.75 s pulsations (Wilms et al.665 2007)., 2007).666 Suzaku performed a TOO observation of the pulsar on 2007 November 30., Suzaku performed a TOO observation of the pulsar on 2007 November 30.667 The results obtained from the analysis of the Suzaku observation are presented in this paper., The results obtained from the analysis of the Suzaku observation are presented in this paper.668 The transient pulsar was observed with the Suzaku during the declining phase of the 2007 November outburst., The transient pulsar was observed with the $Suzaku$ during the declining phase of the 2007 November outburst.669 We used public data (ver-2.1.6.16) for the Suzaku Target of Opportunity (LOO) observation of the pulsar in the present work., We used public data (ver-2.1.6.16) for the $Suzaku$ Target of Opportunity (TOO) observation of the pulsar in the present work.670 The ΝΤΛΟΑΤ monitoring of the pulsar showed that the outburst lasted for 20 davs., The RXTE/ASM monitoring of the pulsar showed that the outburst lasted for $\sim$ 20 days.671 During this outburst. the peak luminosity was about 90 mCrab (7 ASM counts s +).," During this outburst, the peak luminosity was about $\sim$ 90 mCrab $\sim$ 7 ASM counts $^{-1}$ )."672 During this outburst. the pulsar was observed. with the ΗΝΓΙΟ ancl σητα observatories.," During this outburst, the pulsar was observed with the RXTE and $Suzaku$ observatories."673 The RNTE/ASAL (1.5-12 keV) and Swift(BAL (15-50 keV) one-dav averaged light curves of between 2007 September 29 and 2008 January 20 are shown in the top and bottom panels of Figure 1.. respectively.," The RXTE/ASM (1.5-12 keV) and Swift/BAT (15-50 keV) one-day averaged light curves of between 2007 September 29 and 2008 January 20 are shown in the top and bottom panels of Figure \ref{asm}, respectively."674 The region between the vertical lines in the figure. indicates the observation of the pulsar with the Suiza., The region between the vertical lines in the figure indicates the observation of the pulsar with the $Suzaku$.675 This TOO observation was carried out at “XLS nominal” pointing position Lor cllective exposures of 42 ks., This TOO observation was carried out at “XIS nominal” pointing position for effective exposures of 42 ks.676" The ALS was operated. with 71/4 window"" option which gives a time resolution of 2 sec. covering a field of view of 17.8 44."," The XIS was operated with “1/4 window” option which gives a time resolution of 2 sec, covering a field of view of $'$ $\times$ $'$ .4."677 Suzaku. the filth Japanese X-ray astronomy. satellite (Alitsucla et al.," $Suzaku$ , the fifth Japanese X-ray astronomy satellite (Mitsuda et al."678 2007). was launched on 2005 July 10.," 2007), was launched on 2005 July 10."679 [t covers 0.2600 keV. οποιον range with the two sets of, It covers 0.2–600 keV energy range with the two sets of680the maenetized aunulus.,the magnetized annulus.681 The domain is au anuulus with radius raugiug frou 0.08 to 0.32 and height 0.0375 units., The domain is an annulus with radius ranging from 0.08 to 0.32 and height 0.0375 units.682 The vertical boundaries are periodic., The vertical boundaries are periodic.683 The iuner aud outer radial boundaries are fixed-value. arranged by preventing time evolution for particles with radius ereater than 0.3 or less than O.1.," The inner and outer radial boundaries are fixed-value, arranged by preventing time evolution for particles with radius greater than 0.3 or less than 0.1."684 The initial deusitv is 1.0 evervwhere., The initial density is 1.0 everywhere.685" A vertical magnetic field is imposed with radial profile D.(r)=Bob,(re). where By=0.0821 aud which eives a magnetized aunulus."," A vertical magnetic field is imposed with radial profile $B_z(r) = B_0686b_p(r)$, where $B_0=0.0824$ and which gives a magnetized annulus."687" The sound speed im the magnetized auuulus was set το c,=0.821. and the internal enerev in the nonanaguetized region adjusted so that the total pressure (thermal plus magnetic) is conustaut."," The sound speed in the magnetized annulus was set to $c_s = 0.824$, and the internal energy in the non-magnetized region adjusted so that the total pressure (thermal plus magnetic) is constant."688 The radial velocity is perturbed with Spatially constant resolutions of A=1/210. À=1/320. and A=1/100 were used corresponding to A=1/9A\LR. A=1λνμαι. Al/lbAxg where Αν=0.0375 is the wavelength of the fastest erowineg AIRT mode at r=O17.," The radial velocity is perturbed with Spatially constant resolutions of $\lambda=1/240$, $\lambda=1/320$, and $\lambda=1/400$ were used corresponding to $\lambda = 1/9 \lambda_\mathrm{MRI}$, $\lambda = 1/12 \lambda_\mathrm{MRI}$, $\lambda = 1/15 \lambda_\mathrm{MRI}$ where $\lambda_\mathrm{MRI}=0.0375$ is the wavelength of the fastest growing MRI mode at $r=0.17$."689 We then measure the erowth of the most unstable uode at r= 0.17., We then measure the growth of the most unstable mode at $r=0.17$ .690 Fieure 13. shows the radial magnetic Ποια configuration achieved at time 0.91 (or 2.13 orbits at ¢= O17) for all three resolutions., Figure \ref{figflocktestslice} shows the radial magnetic field configuration achieved at time $0.94$ (or $2.13$ orbits at $r=0.17$ ) for all three resolutions.691 To calculate he mode amplitude AZ. we use a convolution defined directly on the particles. iustead of eridding aud Fourier ranstormune the data.," To calculate the mode amplitude $M$, we use a convolution defined directly on the particles, instead of gridding and Fourier transforming the data."692 The motivations are the same as when this procedure was used for the I&elviu-Iehuholtz est., The motivations are the same as when this procedure was used for the Kelvin-Helmholtz test.693 Iu this case. where all suis run over the I points; aud The chosen width of the couvolution σ=2.7«109 imiunnizes the radial range that influcuces the measurement. while still ¢iving sufficieutly low sample niolse.," In this case, where all sums run over the $N$ points, and The chosen width of the convolution $\sigma=2.7\times10^{-6}$ minimizes the radial range that influences the measurement, while still giving sufficiently low sampling noise."694 We plot the evolution of the mode amplitude iu Figure ll together with the maxuuuu erowth rate of exp(0.750) predicted by a linear perturbation analysis of vertical field AMIRI., We plot the evolution of the mode amplitude in Figure \ref{figflocktest_18} together with the maximum growth rate of $\exp(0.75\Omega)$ predicted by a linear perturbation analysis of vertical field MRI.695 The modeled growth rates are reasonably cousisteut with the prediction from tle linear analysis for the fastest erowing mode., The modeled growth rates are reasonably consistent with the prediction from the linear analysis for the fastest growing mode.696 huportautlv. iu the context of the findines of Flocketal.(2010).. where spuriously high erowth rates were observed. we fud erowtli rates slieltly lower than the theoretical maxima value.," Importantly, in the context of the findings of \cite{2010A&A...516A..26F}, where spuriously high growth rates were observed, we find growth rates slightly lower than the theoretical maximum value."697 To determine if Phurbas can be used for practical conrputations. we need to establish some guidelines for its relative ability to resolve particular phenomena.," To determine if Phurbas can be used for practical computations, we need to establish some guidelines for its relative ability to resolve particular phenomena."698 This should be doue cautiously. as different classes of algoritlins have differcut properties iu cach flow regime.," This should be done cautiously, as different classes of algorithms have different properties in each flow regime."699 An equivalence or difference between algorithms in one reenne may not hold across differcut regimes., An equivalence or difference between algorithms in one regime may not hold across different regimes.700 Iu any case. it is expected that an wustructured mesh or unstructured meshless method will have a lower effective resolution than a structured mesh inethod.," In any case, it is expected that an unstructured mesh or unstructured meshless method will have a lower effective resolution than a structured mesh method."701 This is because a given munber of resolving clements can represcut the largest possible set of waveleugths when they are arranged in a regular manner., This is because a given number of resolving elements can represent the largest possible set of wavelengths when they are arranged in a regular manner.702 Civen these caveats. we can compare the test results that we have preseuted here to examples computed with Eulerian. uesh-based schemes.," Given these caveats, we can compare the test results that we have presented here to examples computed with Eulerian, mesh-based schemes."703 The first example is the circularly polarized Alfvéóuu wave test., The first example is the circularly polarized Alfvénn wave test.704 The lowest resolution. three-dimensional. Athena results in a rectaugulhuw domain published in Stoneetal.(2008.Fie.33) correspond to 20 aud 39 cells per waveleugth. computed with third order spatial reconstruction aud IILLD fuxes.," The lowest resolution, three-dimensional, Athena results in a rectangular domain published in \citet[][Fig. 33]{2008ApJS..178..137S}705 correspond to $20$ and $39$ cells per wavelength, computed with third order spatial reconstruction and HLLD fluxes."706 The Phurbas results ou this test at A=1/8 and A=1/16 of a wavceleusth appear to roughly equal the accuracy. of the 20 and 39 cells per wavelength Athena results in the seuse that the final amplitude of the wave in the Phurbas results is closer to the analytically correct value even though there are fewer resolution clemeuts used per waveleneth., The Phurbas results on this test at $\lambda = 1/8$ and $\lambda = 1/16$ of a wavelength appear to roughly equal the accuracy of the $20$ and $39$ cells per wavelength Athena results in the sense that the final amplitude of the wave in the Phurbas results is closer to the analytically correct value even though there are fewer resolution elements used per wavelength.707 This result should be interpreted cautiously. as the two codes lave different aud unoulinear nuuerical dissipation.," This result should be interpreted cautiously, as the two codes have different and nonlinear numerical dissipation."708 Nevertheless. this cau be interpreted to mean that the Phurbas effective resolution A is roughly equal to two Athena cells as a measure of resolution. (," Nevertheless, this can be interpreted to mean that the Phurbas effective resolution $\lambda$ is roughly equal to two Athena cells as a measure of resolution. ("709On average. there should be one particle in cach volume of radius A.),"On average, there should be one particle in each volume of radius $\lambda$ .)"710 Tn this sense Phurbas with a third-order polvnomual fit is competitive with a spatially third-order exid codo., In this sense Phurbas with a third-order polynomial fit is competitive with a spatially third-order grid code.711 A possibly more operationally useful comparison cau be drawn from the results of the linear phase IRI erowtl test., A possibly more operationally useful comparison can be drawn from the results of the linear phase MRI growth test.712 Flockotal.(2010). claim that in the code Pluto (Migeuoueetal.2007) with piecewise linear reconstructions and an HILED σαι solver. 10 cells per wavelength are required to resolve the erowth of AMIRI.," \citet{2010A&A...516A..26F} claim that in the code Pluto \citep{2007ApJS..170..228M} with piecewise linear reconstructions and an HLLD Riemann solver, 10 cells per wavelength are required to resolve the growth of MRI."713 Our test iu section 3.5. shows Phurbas requires ~9 A per MRI wavelength to resolve the exowth., Our test in section \ref{sec_flocktest} shows Phurbas requires $\sim 9$ $\lambda$ per MRI wavelength to resolve the growth.714 Thus. for this test we can sav that one =LA.," Thus, for this test we can say that one $\ \approx 1 \lambda$."715 The alegorithiu used by Flocketal.(2010). has a stencil size of five cells. while Phurbas can be said to have a stencil size of 2ry=LGA. so the same rough proportionality holds between the two algorithms when expressed in terms of stencil size.," The algorithm used by \citet{2010A&A...516A..26F} has a stencil size of five cells, while Phurbas can be said to have a stencil size of $2 r_f =4.6\lambda$, so the same rough proportionality holds between the two algorithms when expressed in terms of stencil size."716 The iain advantage of Plurbas is its Lagrangian mature., The main advantage of Phurbas is its Lagrangian nature.717 Euleriau codes suffer from nunerical dissipation hat varies with the speed aud direction of the flow across the erid., Eulerian codes suffer from numerical dissipation that varies with the speed and direction of the flow across the grid.718 Phurbas formulation cleanly avoids this jehavior., Phurbas's formulation cleanly avoids this behavior.719 For svstenis wheres the bulk velocity varies as a iuultiple or large fraction of the wave speeds. this ueaus Phurbas can capture the dow with more uniformi fidelity across the domain.," For systems where the bulk velocity varies as a multiple or large fraction of the wave speeds, this means Phurbas can capture the flow with more uniform fidelity across the domain."720 Techniques that add an extra advection step to an Eulerian method (e.c.Masset2000: can oulyefficicutly handle simple How geometries.," Techniques that add an extra advection step to an Eulerian method \citep[e.g.][]{2000A&AS..141..165M,2009ApJ...697.1269J} can onlyefficiently handle simple flow geometries."721 Moreover. in Phurbas. the time step is oulv dependent ou Calilean-invariant quantities.," Moreover, in Phurbas, the time step is only dependent on Galilean-invariant quantities."722 For flow with bulk velocities greater than the signal speeds that, For flow with bulk velocities greater than the signal speeds that723as CAIRSs directly. stellar population svnthesis model mus oe considered.,"as CMRs directly, stellar population synthesis model must be considered."724 We adopt a model by Ixodama Arimoto (1997) as such a population svnthesis model., We adopt a model by Kodama Arimoto (1997) as such a population synthesis model.725 Luminosity ancl colour of stars cannot be theoretically estimate without uncertainties because the stellar evolution nioclels hemselves include the uncertainties., Luminosity and colour of stars cannot be theoretically estimated without uncertainties because the stellar evolution models themselves include the uncertainties.726 For example. as for the convection of stellar gas. we have only a phenomenologica heory Cmixing length. theory).," For example, as for the convection of stellar gas, we have only a phenomenological theory (`mixing length theory')."727 However. their model. is adequate for our purpose of investigating the slope of the CAIR.," However, their model is adequate for our purpose of investigating the slope of the CMR."728 Once we understand the properties of galaxies such as the CMIR physically ancl qualitatively. even if the stellar population model is changed. we fit immediately our results with observations by changing parameters mentioned above.," Once we understand the properties of galaxies such as the CMR physically and qualitatively, even if the stellar population model is changed, we fit immediately our results with observations by changing parameters mentioned above."729 The EME which we adopt is the Salpeter type with a slope of 1.35. and the mass range is OAL.~ GOAL..," The IMF which we adopt is the Salpeter type with a slope of 1.35, and the mass range is $0.1$ $_{\odot}\sim 60$ $_{\odot}$."730 The range of stellar metallicity Z; of simple stellar populations is 0.0001~0.05., The range of stellar metallicity $Z_{*}$ of simple stellar populations is $0.0001\sim 0.05$.731 In Section ??.. we divide the stellar component into discl ancl bulge components.," In Section \ref{sec:sf}, we divide the stellar component into disc and bulge components."732 Morphology of cach galaxy is determined. by the £-bancl bulge-to-cise luminosity ratio (B/D)., Morphology of each galaxy is determined by the $B$ -band bulge-to-disc luminosity ratio $B/D$ ).733 Simien de Vaueouleurs (1986). showed. that he Llubble tvpe of galaxies correlates with the D-band uminosity 2/0., Simien de Vaucouleurs (1986) showed that the Hubble type of galaxies correlates with the $B$ -band luminosity $B/D$ .734 In this paper. galaxies with 2/251.52 are identified as ellipticals. 0.680<1.52 as SOs. and DfD<0.68 as spirals. according to their results.," In this paper, galaxies with $B/D\geq 1.52$ are identified as ellipticals, $0.68\leq B/D<1.52$ as S0s, and $B/D<0.68$ as spirals, according to their results."735 Lt is shown hat this method for classification reproduces observations well by Ixaullmann et al. (, It is shown that this method for classification reproduces observations well by Kauffmann et al. (7361993) and Baugh. Cole brenk (1996).,"1993) and Baugh, Cole Frenk (1996)."737 In this section. we explore the origin of the CMI," In this section, we explore the origin of the CMR."738 In this paper. because we investigate how the CALR depends on the physical processes such as star formation. supernova fecdback. and so on. we fix the cosmological parameters to the standard CDM model. that is. £2=1A0. 50km | +. ο=0.06. and os=0.67.," In this paper, because we investigate how the CMR depends on the physical processes such as star formation, supernova feedback, and so on, we fix the cosmological parameters to the standard CDM model, that is, $\Omega=1,739\Lambda=0, H=50$ km $^{-1}$ $^{-1}$, $\Omega_{b}=0.06$, and $740\sigma_{8}=0.67$."741 We refer to the moclels considered bv the models A. D. € aad D as shown in Table 1.," We refer to the models considered by the models A, B, C and D as shown in Table 1."742 In this table. we also show models from E to J. which are adopted in the next section.," In this table, we also show models from E to J, which are adopted in the next section."743 The fifth column. UV. is explained in Section ??..," The fifth column, UV, is explained in Section \ref{sec:uv}."744 Phe sixth column. burst. is shown in Section ??..," The sixth column, burst, is shown in Section \ref{sec:merger}."745 The vield y is equal to 0.038=2Z. in all the models., The yield $y$ is equal to $0.038=2Z_{\odot}$ in all the models.746 In WO. y=1.22. in the low feedback model.," In KC, $y=1.2Z_{\odot}$ in the low feedback model."747 Lowever. in order to see the ellects of the feedback and so on. we fix the value of the vield in all models.," However, in order to see the effects of the feedback and so on, we fix the value of the yield in all models."748 1n the following sections. we take notice of only the slopes of CMIti. because the luminosity of galaxies can be translated by considering the following reason.," In the following sections, we take notice of only the slopes of CMRs, because the luminosity of galaxies can be translated by considering the following reason."749 Stars are formed. according to the IME., Stars are formed according to the IMF.750 Mass of luminous stars is larger than O.OSAL.. which is determined by the criterion of nuclear burning.," Mass of luminous stars is larger than $\sim 0.08$ $_{\odot}$, which is determined by the criterion of nuclear burning."751 However. there is a possibility hat invisible stars with mass smaller than O.OSAL. are ormed.," However, there is a possibility that invisible stars with mass smaller than $0.08$ $_{\odot}$ are formed."752 LW there are many invisible stars in galaxies. the galaxies become faint compared to the case that all stars are uminous.," If there are many invisible stars in galaxies, the galaxies become faint compared to the case that all stars are luminous."753 “Phe ratio of the invisible stars to the luminous stars is (treated: as a [ree parameter in the previous work., The ratio of the invisible stars to the luminous stars is treated as a free parameter in the previous work.754 Therefore the absolute value of luminosity can be adjusted w this parameter., Therefore the absolute value of luminosity can be adjusted by this parameter.755 Lhe suitable value of this parameter will » determined by considering other observational quantities., The suitable value of this parameter will be determined by considering other observational quantities.756 Thus this parameter does not alfect the slope of the CAI., Thus this parameter does not affect the slope of the CMR.757 1n . we show the CMldIis in the models A and D with Aner=2 which is the value adopted by WC.," In \ref{fig:cmrab}, we show the CMRs in the models A and B with $\alpha_{hot}=2$ which is the value adopted by KC."758 Ehe dots denote galaxies identified as ellipticals. and the solid lines show the observational CARs (Bower οἱ al.," The dots denote galaxies identified as ellipticals, and the solid lines show the observational CMRs (Bower et al."759 1992)., 1992).760 The criterion to xck out ellipticals among all galaxies is shown in Section 7aed., The criterion to pick out ellipticals among all galaxies is shown in Section \ref{sec:mor}.761 Asin KC. in the model A with high feedback clliciency. he slopes of the CAIRs are in roughlv agreement with he observations. but the dispersion is larger than that of observations. 0.04 mag.," As in KC, in the model A with high feedback efficiency, the slopes of the CMRs are in roughly agreement with the observations, but the dispersion is larger than that of observations, $\sim 0.04$ mag."762 This is due to the recent. star ormation owing to long star formation time-scale. τι=20 Gyr.," This is due to the recent star formation owing to long star formation time-scale, $\tau_{*}^{0}=20$ Gyr."763 On the other hand. in the model B with low feedback elliciency. the slopes of the οΑς are nearly Lat.," On the other hand, in the model B with low feedback efficiency, the slopes of the CMRs are nearly flat."764 At the xieht-end. the colour becomes redder about O.1-0.2 mag. and at the faint-enc (My~ 15). about. 0.2-0.3 mag. compared to that in the modelA. In order to see the physical relation between the reddening and the feedback intensity. we investigate the age- and metallicitv-Iuminosity relations.," At the bright-end, the colour becomes redder about 0.1-0.2 mag, and at the faint-end $M_{V}\sim -18$ ), about 0.2-0.3 mag, compared to that in the modelA. In order to see the physical relation between the reddening and the feedback intensity, we investigate the age- and metallicity-luminosity relations."765Fig.,Fig.766 10 presents contour plots of stars in the simulated rolwine ealaxy at f=1600 Myrs (820 Myrs after the oerieenter passage)., 10 presents contour plots of stars in the simulated polar-ring galaxy at $t=1600$ Myrs (820 Myrs after the pericenter passage).767 At that time the donor galaxy has flow away at about 250 kpe distance (which of course. can appear much closer in sky projection) and does nof erturb the target anv more.," At that time the donor galaxy has flown away at about 250 kpc distance (which of course, can appear much closer in sky projection) and does not perturb the target any more."768 As for the polar ring. it has ecole more regular than on the earlier stages shown iu Fig.," As for the polar ring, it has become more regular than on the earlier stages shown in Fig."769 9. even if still asviinuetrical.," 9, even if still asymmetrical."770 This model reproduces he main characteristics of AM. 1931-563: Let us note that the matter transferred from the douor o the target galaxy is not only gas., This model reproduces the main characteristics of AM 1934-563: Let us note that the matter transferred from the donor to the target galaxy is not only gas.771 As can be seen iu Fig., As can be seen in Fig.772 9. some stars frou the donor are also captured i the xolar disk. without being dispersed. but most of them have )cen formed just before the ring. inside the tidal bridgc.," 9, some stars from the donor are also captured in the polar disk, without being dispersed, but most of them have been formed just before the ring, inside the tidal bridge."773 That soue stars are formed before the ring mav lead to an over-estiniatioi of the age of the ring. vet these stars do rot dominate the mass and are not much older than the ring. so that this over-cstimation may nof excece 24107 vr.," That some stars are formed before the ring may lead to an over-estimation of the age of the ring, yet these stars do not dominate the mass and are not much older than the ring, so that this over-estimation may not exceed $\times 10^8$ yr."774" A sinall fraction of dark matter is also acereted frou hne donunor eaaxy. but for our model with non-rotatiug dark haoes, this does not exceed of the domor dark Lass,"," A small fraction of dark matter is also accreted from the donnor galaxy, but for our model with non-rotating dark haloes, this does not exceed of the donnor dark mass."775 At least four scenarios cau be proposed to account for voung. massive polar rugs surrounding pre-existing host ealaxies: The last three ones cau be regarded as hree subtypes the accretion scenark»," At least four scenarios can be proposed to account for young, massive polar rings surrounding pre-existing host galaxies: The last three ones can be regarded as three subtypes of the accretion scenario."776 Let Us note that the first and third of these variants do no predic the presence : conirpanionus around PRCs. wuch is supported by servations.," Let us note that the first and third of these variants do not predict the presence of companions around PRGs, which is supported by observations."777 For example Brocea et al. (," For example, Brocca et al. ("7781997) find that statistically. the environments of PRCs show no excess of ose coniaΠοis. which is consisteit with the uajoritv of PRCs forming via either kong-ternu sectlay gas accretion or via mergers In which the οςΜΙΣΟΙ Is estrovec.,"1997) find that statistically, the environments of PRGs show no excess of close companions, which is consistent with the majority of PRGs forming via either long-term secular gas accretion or via mergers in which the companion is destroyed."779 Towever. this does not 11le out accretion Yolla eas-ricli donor. for the donor may have left he scene.," However, this does not rule out accretion from a gas-rich donor, for the donor may have left the scene."780 Torice et: (2002a.b) favor the merger scenark> of two ¢isk. ealaxies. frou the study of stelar »opulatious and morphological structure of the host: however. the costraints are 1O strong euouel. Since du oan scenario. enas can be acCreec by t1ο host aud mduce star formation. wlie the host Is porurbed. so that a large range of 5‘lar ages sLOU.d coexist in the rturbed host.," Iodice et al (2002a,b) favor the merger scenario of two disk galaxies, from the study of stellar populations and morphological structure of the host; however, the constraints are not strong enough, since in any scenario, gas can be accreted by the host and induce star formation, while the host is perturbed, so that a large range of stellar ages should coexist in the perturbed host."781" The prex""ce or Lot of a diffuse eivelope of stars arouud the PRC is LOVE COISune (see BC03).", The presence or not of a diffuse envelope of stars around the PRG is more constraining (see BC03).782 The major merger «xnario las 1 shown to be ess likely than the fida accretion iulo. especially or inclined rings {06000).," The major merger scenario has been shown to be less likely than the tidal accretion scenario, especially for inclined rings (BC03)."783 Aloreover. 10 asviunetrie rue of AM. 1931-563 cinuuot νο well fitted bv. this scenario.," Moreover, the asymmetric ring of AM 1934-563 cannot be well fitted by this scenario."784 Tudeed. this scenido assunes that the host σαιAXV as ποσο with a “victim cisk galaxy iat has even bir ο the polar ring.," Indeed, this scenario assumes that the host galaxy as merged with a ”victim” disk galaxy that has given birth to the polar ring."785" Since t1ο southern part of the riie ] ANL 1931-563 is more exteided than t16 northern onc. f center of this “victim, σαaxv should o south of the host ealaxy. which should ial: ot1e southern part of the volar ving brieliter that t16 OLheru one."," Since the southern part of the ring in AM 1934-563 is more extended than the northern one, the center of this “victim” galaxy should be south of the host galaxy, which should make the southern part of the polar ring brighter that the northern one."786 Ou the contrary. f1ο rorthern part o the ving is brighter than the southeru one.," On the contrary, the northern part of the ring is brighter than the southern one."787 We have cjecked. iu ESuulatious of DCUS that t1C nore extended AL Q: the ring is at least as hunuinous astje shortest part. when tie Ting is formed in a major nerser.," We have checked in simulations of BC03 that the more extended part of the ring is at least as luminous as the shortest part, when the ring is formed in a major merger."788 We lave explaiuec aboxο that the disruption of a small companion. duriug a niünor merecr. is uulikelv to lave orlned the polar ring of AM 1931-563. for the rine NOTE certainly appear unuclosed. because the orbit of the conauion is uuclosed iself.," We have explained above that the disruption of a small companion, during a minor merger, is unlikely to have formed the polar ring of AM 1934-563, for the ring would certainly appear unclosed, because the orbit of the companion is unclosed itself."789 Moreover. forming a massive enough ring. without disturbing the host disk too wich. In asrong constraint on the mass of this companion.," Moreover, forming a massive enough ring, without disturbing the host disk too much, is a strong constraint on the mass of this companion."790 ο1 the contrary. f1ο accretion scenario invoking a fida nass transfer from a massive donor. without galaxy merger. has been shown to succeed in reproducing the characteristics of AM |931-563.," On the contrary, the accretion scenario invoking a tidal mass transfer from a massive donor, without galaxy merger, has been shown to succeed in reproducing the characteristics of AM 1934-563."791 Furthermore. the small velocity dispersion of this group (see Sect.," Furthermore, the small velocity dispersion of this group (see Sect."792 3.1) could result, 3.1) could result793produce additional. structure at a few given. Lrequencics.,produce additional structure at a few given frequencies.794 The general scheme of propagating Iluctuations can still be correct but. additional assumptions need to be mace. for example. assuming that a few annuli have much enhanced variability power will certainly. produce a bumpy PSD and might reproduce the steps in the lag spectra. as suggested by Nowak(2000). and also Ixotovetal.(2001).," The general scheme of propagating fluctuations can still be correct but additional assumptions need to be made, for example, assuming that a few annuli have much enhanced variability power will certainly produce a bumpy PSD and might reproduce the steps in the lag spectra, as suggested by \citet{Nowak00} and also \citet{Kotov}."795. The extended. emitting region in our model is responsible for the high frequency bend in the PSD. and the associated racial emissivity profiles produce the energy dependence of the PSD and the time lags.," The extended emitting region in our model is responsible for the high frequency bend in the PSD, and the associated radial emissivity profiles produce the energy dependence of the PSD and the time lags."796 In this section we will summarise the speetral-timing properties produced. by the. moclel. consider the implications of the model for the size of the emitting region. and discuss some possible improvemoents.," In this section we will summarise the spectral-timing properties produced by the model, consider the implications of the model for the size of the emitting region, and discuss some possible improvements."797 Keeping the local variability time-scales tied. το the propagation time-scale produces time lage spectra of power law slope ~l or Uatter. and lags of ~110% of the variability time-scale. for a wide range of moclel parameters.," Keeping the local variability time-scales tied to the propagation time-scale produces time lag spectra of power law slope $\sim -1$ or flatter, and lags of $\sim 1 -10\%$ of the variability time-scale, for a wide range of model parameters."798 As discussed in Sections 4. and 5.. these simple assumptions produce lag spectra that mateh the cata well in time-scale dependence ancl amplitude.," As discussed in Sections \ref{agn}799 and \ref{xrb}, these simple assumptions produce lag spectra that match the data well in time-scale dependence and amplitude."800 We note once again that these laes arise solely due to the clillerence in emissivity. profiles of the N-rav. energy. bands ancl do not involve any other spectral evolution of the emitting region., We note once again that these lags arise solely due to the difference in emissivity profiles of the X-ray energy bands and do not involve any other spectral evolution of the emitting region.801 The amplitude of the lags depends mostly on the emissivity indices of the energy. bands. increasing rapiclly with their difference. up to As~1. above which the lag values tend to saturate.," The amplitude of the lags depends mostly on the emissivity indices of the energy bands, increasing rapidly with their difference, up to $\Delta \gamma \sim 1$, above which the lag values tend to saturate."802 Significant lags can appear between enerev bands. characterised by similar emissivity profiles and. correspondingly. similar PSD shapes.," Significant lags can appear between energy bands characterised by similar emissivity profiles and, correspondingly, similar PSD shapes."803 In particular. the PSD of €vg X-1 in the high/soft state shows weak energy dependence (see Section 5.1). but the PSD ratio and lags can still be reproduced simultaneously with close emissivity indices for each energy band (As~ 0.5). provided that the propagation time-scale is slightly longer than the local luctuation time-scale. and assuming that the inner clise radius is small.," In particular, the PSD of Cyg X-1 in the high/soft state shows weak energy dependence (see Section 5.1), but the PSD ratio and lags can still be reproduced simultaneously with close emissivity indices for each energy band $\Delta \gamma \sim 0.5$ ), provided that the propagation time-scale is slightly longer than the local fluctuation time-scale, and assuming that the inner disc radius is small."804 Ehe power-law shape of the lag spectra is quite robust. its slope depends. only weakly on. disc structure parameters (LfB)-a.," The power-law shape of the lag spectra is quite robust, its slope depends only weakly on disc structure parameters $(H/R)^2\alpha$."805 Incidentally. the stability of the lag spectra might explain the behaviour seen in. e.g. (νο X-1. in different spectral states.," Incidentally, the stability of the lag spectra might explain the behaviour seen in, e.g. Cyg X-1, in different spectral states."806 Fhis object shows very dillerent PSD and energy spectra in the high/soft and Iowhard. states. indicative of dillerent. disc configurations. but surprisingly similar lags (Pottschmidtetal.," This object shows very different PSD and energy spectra in the high/soft and low/hard states, indicative of different disc configurations, but surprisingly similar lags \citep{Pottschmidt}."8072000).. A comparison with AGN and BUNKB data in the high/soft state shows that the filtering ellect of the extended emitting region. acting on a simple 1/f underlying PSD shape. can broadly reproduce their PSD shape and energy dependence.," A comparison with AGN and BHXRB data in the high/soft state shows that the filtering effect of the extended emitting region, acting on a simple $1/f$ underlying PSD shape, can broadly reproduce their PSD shape and energy dependence."808 Cve N-1 in the low/hare state. however. requires a more complex PSD.," Cyg X-1 in the low/hard state, however, requires a more complex PSD."809 In either case. the extended emitting region introduces a bend in the PSD in addition to any intrinsic curvature. ancl produces the energy. dependence of he filtered. PSDs.," In either case, the extended emitting region introduces a bend in the PSD in addition to any intrinsic curvature, and produces the energy dependence of the filtered PSDs."810 Note that the bending power-Iaw. used to fit AGN data. is only an approximation to the actual PSDs produced. by he fluctuating-acerction model.," Note that the bending power-law, used to fit AGN data, is only an approximation to the actual PSDs produced by the fluctuating-accretion model."811 The filtered. PSD bends down continuously at. high. frequencies ancl has no well-defined. high-frequency power Law slope., The filtered PSD bends down continuously at high frequencies and has no well-defined high-frequency power law slope.812 The cdillerence tween the bending power-law mocel and the filtered PSDs can be appreciated in Fig. Ll.," The difference between the bending power-law model and the filtered PSDs can be appreciated in Fig. \ref{psd_surr_ngc},"813 where the single-bend PSD it overestimates the power in the highest. frequency. bins., where the single-bend PSD fit overestimates the power in the highest frequency bins.814 Llowever. as this region of the PSDs from real data is often wavily alected by Poisson noise. the associated. error bars and scatter are large and it is not. possible to appreciate any deviations from a simple power law slope.," However, as this region of the PSDs from real data is often heavily affected by Poisson noise, the associated error bars and scatter are large and it is not possible to appreciate any deviations from a simple power law slope."815 Therefore. even better ACN data would. be needed to discern if the model can replicate the exact. PSD shape or if additional variability components are needed.," Therefore, even better AGN data would be needed to discern if the model can replicate the exact PSD shape or if additional variability components are needed."816 We also note here that. although Itevnivtsev.Gilfanov&Churazov(2000). fit the ügh/soft state PSD from combined data from 1996 June 4-18 with a simple power law (index -2.1) up to ~200 112. he signal-to-noise at these frequencies is still fairly low. and he PSD during that time is known vary in shape between observations (Cuietal.1997)... which makes interpretation of the underlving shape even more dillicult.," We also note here that, although \citet{RevHifreq} fit the high/soft state PSD from combined data from 1996 June 4-18 with a simple power law (index -2.1) up to $\sim200$ Hz, the signal-to-noise at these frequencies is still fairly low, and the PSD during that time is known vary in shape between observations \citep{Cui97}, which makes interpretation of the underlying shape even more difficult."817" Axclssonctal.(2005) have recently demonstrated that the high/soft state ""SD can be fitted with weak Lorentzians to a ~Lf power-law component with an exponential cut-olf.", \citet{Axelsson} have recently demonstrated that the high/soft state PSD can be fitted with weak Lorentzians to a $\sim 1/f$ power-law component with an exponential cut-off.818 This last component is reminiscent of our simulated PSDs., This last component is reminiscent of our simulated PSDs.819 Therefore. our model may be considered as representative of the times when the Lorentzians in the PSD are verv weak or absent.," Therefore, our model may be considered as representative of the times when the Lorentzians in the PSD are very weak or absent."820 In our model. thick cise parameters. a(44/R)?=0.3. in the inner regions of the accretion Low and inner radius Moin=6 put the break frequency. around. 10eIi. (e. 2 - 20 Hz for a LOAD. black hole. 103 for a 10A4. ).," In our model, thick disc parameters, $\alpha (H/R)^2=0.3$, in the inner regions of the accretion flow and inner radius $r_{\rm min}= 6$ put the break frequency around $10^{-4}-10^{-3}c/R_{\rm g}$ (i.e. $\sim$ 2 - 20 Hz for a $10 M_{\odot}$ black hole, $2\times82110^{-5}-2\times 10^{-4}$ for a $10^6 M_{\odot}$ )."822" ""οσο values are in general2. agreement with the average break frequencies of ΑΝ and. DIIXIUDs. indicating that Ductuations on the viscous time-scales of a gcometrically Chick accretion Low are appropriate to explain he variability."," These values are in general agreement with the average break frequencies of AGN and BHXRBs, indicating that fluctuations on the viscous time-scales of a geometrically thick accretion flow are appropriate to explain the variability."823 A geometrically thick disc is necessary to srocluce the observed high frequency Huctuations as frequencies., A geometrically thick disc is necessary to produce the observed high frequency fluctuations as frequencies.824 I£ the [uctuations were instead. produced on dynamical time-scales. or à magnetic time-scale related ο this. as in e.g. Wineetal.(2004).. then a thinner Low might be allowed.," If the fluctuations were instead produced on dynamical time-scales, or a magnetic time-scale related to this, as in e.g. \citet{King}, then a thinner flow might be allowed."825 However. the Uuctuations need. not only »* produced: but also propagated. which poses a dilliculty or much thinner accretion Lows. as fluctuations on time-scales much shorter than the propagation time-scale are casily clampecl (Churazoyetal.," However, the fluctuations need not only be produced but also propagated, which poses a difficulty for much thinner accretion flows, as fluctuations on time-scales much shorter than the propagation time-scale are easily damped \citep{Churazov}."8262001)... This. of course. does not rule out an additional thin disc. possibly underlving the thick How. that might contribute to the Dux but not to the variability.," This, of course, does not rule out an additional thin disc, possibly underlying the thick flow, that might contribute to the flux but not to the variability."827 In our implementation of the variability model. the X-ray emitting region extends out to large radii.," In our implementation of the variability model, the X-ray emitting region extends out to large radii."828 Phe bend in the PSD is produced. by the radial distribution of variability time-scales anc cniissivity profiles alone and is not related to a characteristic time-scale at the maximum racius of emission., The bend in the PSD is produced by the radial distribution of variability time-scales and emissivity profiles alone and is not related to a characteristic time-scale at the maximum radius of emission.829 l'herefore. an outer οσο to the emitting region might be allowed. but is not required.," Therefore, an outer edge to the emitting region might be allowed but is not required."830 The steep emissivity, The steep emissivity831The profile of angular velocity in (he convection zone is determined by a balance of angular momentum transport [rom meridional flow and a reduction in meridional How from buovaney force at the subacliabatic laver.,The profile of angular velocity in the convection zone is determined by a balance of angular momentum transport from meridional flow and a reduction in meridional flow from buoyancy force at the subadiabatic layer.832" We run simulations for seventeen cases. with Table | showing the parameters for each Case,"," We run simulations for seventeen cases, with Table \ref{param} showing the parameters for each case."833 In this section. we discuss (he cases wilh angular velocities up to 16 times the solar value (represented by Q.). placing an emphasis on the morphology. of stellar differential rotation.," In this section, we discuss the cases with angular velocities up to 16 times the solar value (represented by $\Omega_\odot$ ), placing an emphasis on the morphology of stellar differential rotation."834 Fie., Fig.835 3 shows the results of our calculations which correspond to cases 1-5 in Table 1.., \ref{rapid} shows the results of our calculations which correspond to cases 1-5 in Table \ref{param}.836 It is found that the larger stellar angular velocity is. the more likely it is [or differential rotation to be in the Tavlor-Proudiman state. in whieh the contour lines of the," It is found that the larger stellar angular velocity is, the more likely it is for differential rotation to be in the Taylor-Proudman state, in which the contour lines of the"837implies a nean iron abundance of about 0.25 dex at 5 kpe towards the inner disk (and -0.25 towards the outer disk).,implies a mean iron abundance of about 0.25 dex at 5 kpc towards the inner disk (and -0.25 towards the outer disk).838 Allowing for an intrinsic dispersion similar to what is measured on the local iron distribution (0.15 dex). we easily reach (within 1 or 2 sigma) the highest metallicities that are observed at the solar radius. which means that the oxvgen gradient measured on Type II PNe is well compatible with abundances measured in the local stellar population.," Allowing for an intrinsic dispersion similar to what is measured on the local iron distribution (0.15 dex), we easily reach (within 1 or 2 sigma) the highest metallicities that are observed at the solar radius, which means that the oxygen gradient measured on Type II PNe is well compatible with abundances measured in the local stellar population."839 In order to increase (he significance of the PN results we explore the literature Lor object classes whose a-element abundances and gradients have been studied., In order to increase the significance of the PN results we explore the literature for object classes whose $\alpha$ -element abundances and gradients have been studied.840 Several studies concerning II IE regions and voung stars are available. while the only possible comparison with an older population is that of open clusters.," Several studies concerning H II regions and young stars are available, while the only possible comparison with an older population is that of open clusters."841 In Table 4 we list gradient slopes and (heir uncertainties (in column 2) and distance ranges (column 3) from different (racers (column 4). where the authors have measured directly the oxveen abundances of the (racers.," In Table 4 we list gradient slopes and their uncertainties (in column 2) and distance ranges (column 3) from different tracers (column 4), where the authors have measured directly the oxygen abundances of the tracers."842 Most tracers are very voung stars. whose eradients have been plotted in Figure 8 with open circles.," Most tracers are very young stars, whose gradients have been plotted in Figure 8 with open circles."843 These voung stellar data seem {ο fit in well with the voung PNe of our sample. eiving continuity to the plot.," These young stellar data seem to fit in well with the young PNe of our sample, giving continuity to the plot."844 Our basic list of open clusters is (he one designed by Magrini et al. (, Our basic list of open clusters is the one designed by Magrini et al. (8452009). to which we added a few more clusters from (he most recent literature.,"2009), to which we added a few more clusters from the most recent literature."846 The various parameters are eiven in Table 5 with columns (1) through (2) providing the name. age. distance from the Galactic center.," The various parameters are given in Table 5 with columns (1) through (3) providing the name, age, distance from the Galactic center."847 Iron abundances. referred. to solar. are listed in column (4).," Iron abundances, referred to solar, are listed in column (4)."848 References to the iron abundances are as in Magrini et al. (, References to the iron abundances are as in Magrini et al. (8492009).,2009).850 Columns (5) through (8) of Table 5 give the oxvgen abundance ratios (o iron. the oxvgen abundances relative {ο solar. the actual oxvgen abundance caleulated for the solar value given in the individual papers. anc," Columns (5) through (8) of Table 5 give the oxygen abundance ratios to iron, the oxygen abundances relative to solar, the actual oxygen abundance calculated for the solar value given in the individual papers, and"851 Columns (5) through (8) of Table 5 give the oxvgen abundance ratios (o iron. the oxvgen abundances relative {ο solar. the actual oxvgen abundance caleulated for the solar value given in the individual papers. ancl," Columns (5) through (8) of Table 5 give the oxygen abundance ratios to iron, the oxygen abundances relative to solar, the actual oxygen abundance calculated for the solar value given in the individual papers, and"852"As discussed in3.1,, when the model parameters that determine the density and temperature profiles at TA are unconstrained or poorly constrained, the fully parametric approach can be biased towards self-similar scaling relations.","As discussed in, when the model parameters that determine the density and temperature profiles at $r_\Delta$ are unconstrained or poorly constrained, the fully parametric approach can be biased towards self-similar scaling relations."853" As an explicit illustration of this, consider a 8 model description of the gas density in conjunction with a simple, non-isothermal temperature profile, This function is a simplification of the form used by(2006),, namely eliminating the ‘cool core’ term, which is intended to describe the profile at small radii."," As an explicit illustration of this, consider a $\beta$ model description of the gas density in conjunction with a simple, non-isothermal temperature profile, This function is a simplification of the form used by, namely eliminating the `cool core' term, which is intended to describe the profile at small radii."854" To illustrate the case where these models are effectively unconstrained, we generated random realizations by sampling independent, uniform values of the model parameters within the ranges given inA1."," To illustrate the case where these models are effectively unconstrained, we generated random realizations by sampling independent, uniform values of the model parameters within the ranges given in."855". The radial scales in the density and temperature models, Τε and Τε, were allowed to take values between zero and raa. max4/3f T"," The radial scales in the density and temperature models, $\rc$ and $\rt$, were allowed to take values between zero and $\rmax = \mathrm{max} \sqrt{3\beta T_0}$ ."856"his is the maximum value of rc for which the isothermal To.8 model has a real solution for ra 10)); while the same is not true of this non-isothermal model, allowing larger values does not change the resulting picture qualitatively."," This is the maximum value of $\rc$ for which the isothermal $\beta$ model has a real solution for $r_\Delta$ ); while the same is not true of this non-isothermal model, allowing larger values does not change the resulting picture qualitatively."857" 8 was allowed to vary over a range somewhat wider than that seen in observations, while the temperature exponents, b and c, were varied over approximately the range allowed by(2006)."," $\beta$ was allowed to vary over a range somewhat wider than that seen in observations, while the temperature exponents, $b$ and $c$, were varied over approximately the range allowed by."858". To provide an adequate baseline to observe the resulting scaling behavior, the temperature normalization, To, was sampled uniformly in the logarithm between 1 and 1000."," To provide an adequate baseline to observe the resulting scaling behavior, the temperature normalization, $T_0$, was sampled uniformly in the logarithm between 1 and 1000."859" For each realization, an implicit solution for ra 3)) was searched for numerically, and models for which there was no real solution were discarded."," For each realization, an implicit solution for $r_\Delta$ ) was searched for numerically, and models for which there was no real solution were discarded."860 A sample of the resulting density and temperature profiles is shown inAl., A sample of the resulting density and temperature profiles is shown in.861". The model profiles are clearly not self-similar in any meaningful sense, but, because their variation is independent of mass, the slopes of the mean scaling relations take on the self-similar values (right panel of 1))."," The model profiles are clearly not self-similar in any meaningful sense, but, because their variation is independent of mass, the slopes of the mean scaling relations take on the self-similar values (right panel of )."862" For clarity, we have culled models where rc/rA>0.7 from the figure; these models produce an asymmetric tail to low masses, but do not change the scaling relation slope."," For clarity, we have culled models where $\rc/r_\Delta>0.7$ from the figure; these models produce an asymmetric tail to low masses, but do not change the scaling relation slope."863" The x-axis of the figure shows the emission-weighted projected temperature within ra, calculated from the n(r) and T(r) profiles, although the precise definition of T does not affect the conclusions."," The x-axis of the figure shows the emission-weighted projected temperature within $r_\Delta$, calculated from the $n(r)$ and $T(r)$ profiles, although the precise definition of $T_\Delta$ does not affect the conclusions."864 Here we offer some brief thoughts on the task of obtaining hydrostatic mass estimates using minimal assumptions, Here we offer some brief thoughts on the task of obtaining hydrostatic mass estimates using minimal assumptions865age of the Universe at each redshift considered).,age of the Universe at each redshift considered).866 While there is some room for confusion with z«1 galaxies — which due to dust emission in the F560W filter (not taken into account in our models) may potentially produce similar colours in this diagram — such objects are not likely to display apparent magnitudes in the same range as high-redshift pop III galaxies., While there is some room for confusion with $z<1$ galaxies – which due to dust emission in the F560W filter (not taken into account in our models) may potentially produce similar colours in this diagram – such objects are not likely to display apparent magnitudes in the same range as high-redshift pop III galaxies.867" Hence, objects that end up in the upper left corner of Fig. 8,"," Hence, objects that end up in the upper left corner of Fig. \ref{typeA_MIRIcolcol},"868" and have apparent magnitudes in the range expected for high-redshift galaxies, are likely to be zzz 7-8 pop III galaxies even in the absence of additional redshift constraints."," and have apparent magnitudes in the range expected for high-redshift galaxies, are likely to be $z\approx 7$ –8 pop III galaxies even in the absence of additional redshift constraints."869" While this scheme would seem to give a cleaner selection of pop III galaxy candidates than that proposed by Inoue(2011b),, it suffers from one obvious drawback: the lower sensitivity of MIRI implies a minimum mass for detection that is an order of magnitude higher than in the case where only NIRCam filters are used (see Fig. 5))."," While this scheme would seem to give a cleaner selection of pop III galaxy candidates than that proposed by \citet{Inoue b}, it suffers from one obvious drawback: the lower sensitivity of MIRI implies a minimum mass for detection that is an order of magnitude higher than in the case where only NIRCam filters are used (see Fig. \ref{Mmin_singlez}) )."870" Even in the case of 100 h exposures per filter, the mass converted into stars would need to be on the order of ~10’Mo to allow detection in the F560W and F770W filters."," Even in the case of 100 h exposures per filter, the mass converted into stars would need to be on the order of $\sim 10^7\ M_\odot$ to allow detection in the F560W and F770W filters."871" The maximum mass that stellar populations consisting entirely of pop III stars can reach is unknown, but current simulations suggest that unenriched halos are unlikely to attain total masses in excess of M~105Mg at z>7 (Trentietal. 2009)."," The maximum mass that stellar populations consisting entirely of pop III stars can reach is unknown, but current simulations suggest that unenriched halos are unlikely to attain total masses in excess of $M\sim 10^8\ M_\odot$ at $z>7$ \citep{Trenti et al.}."872". To produce~10°Meg worth of pop III stars, essentially all of the baryons in a M~105Mo object would need to be converted into stars (somehow evading negative feedback effects) in a limited amount of time (up to ~107 yr), which seems highly unrealistic."," To produce$\sim 10^7\ M_\odot$ worth of pop III stars, essentially all of the baryons in a $M\sim 10^8\ M_\odot$ object would need to be converted into stars (somehow evading negative feedback effects) in a limited amount of time (up to $\sim 10^7$ yr), which seems highly unrealistic."873" By hunting for pop III galaxies behind lensing clusters with magnification µzz100 (MACSJ0717.5+3745;e.g.Zitrin 2010),, the required stellar population mass can in principle be lowered to ~10°Mo "," By hunting for pop III galaxies behind lensing clusters with magnification $\mu\approx 100$ \citep[MACS J0717.5+3745; e.g.][]{Zitrin et al.,Zackrisson et al. c}, , the required stellar population mass can in principle be lowered to $\sim 10^5\ M_\odot$ "874[or (10)) are sunimarized in Table 1 for different: phases.,for \ref{107a}) ) are summarized in Table \ref{tab1} for different phases.875 lig., Fig.876 lis a plot of Πίο). 962). log gel). ο) for one whole pulsation. INSGS).," 1 is a plot of $T_{\rm e}(\varphi)$ , $\vartheta(\varphi)$, $\log g_{\rm e}(\varphi)$ , $h_0(R,\varphi)$ for one whole pulsation. $\Delta T_{\rm e}(\varphi)$,"877 Aloeg(y2) are plotted in the lower panels of Fig. 2.., $\Delta\log g_{\rm e}(\varphi)$ are plotted in the lower panels of Fig. \ref{fig2}.878 Assuming random errors of 40.02 for the colour indices will result in AY.=Εθν and Aloeg.=+£0.04., Assuming random errors of $\pm 0.02$ for the colour indices will result in $\Delta T_{\rm e}=\pm 10\mbox{K}$ and $\Delta\log g_{\rm e}=\pm 0.04$.879 These values are indicated by the dotted horizontal lines., These values are indicated by the dotted horizontal lines.880 Condition Lis satisfied in the phase points Lvine below or close to the clotted lines., Condition I is satisfied in the phase points lying below or close to the dotted lines.881 At phases [ving. above the dotted. lines. the monochromatic Hux of statie models and SU Dra dillers significantly on a level which has noticeable ellect on the broad band colours (BV(De:," At phases lying above the dotted lines, the monochromatic flux of static models and SU Dra differs significantly on a level which has noticeable effect on the broad band colours $UBV(RI)_C$."882 The assumed Al] L6.E(BVW)=0.015 were verified by a variation procedure (Bareza&Benkó2009).," The assumed $[M]=-1.6$, $E(B-V)=0.015$ were verified by a variation procedure \citep{barc3}."883. In the shock free phases yo=0.15.0.5.0.55. minimization of AT.(uz).Adoggol) resulted in A]=100-010. E(BV)=0.015£0.01.," In the shock free phases $\varphi=0.15,0.5,0.55$, minimization of $\Delta884T_{\rm e}(\varphi),\Delta\log g_{\rm e}(\varphi)$ resulted in $[M]=-1.60\pm 0.10$, $E(B-V)=0.015\pm 0.01$."885 To demonstrate the difference between good ancl poor QSAA. the histograms of the 30 logge.Ze values are plotted in the upper panels of Fig.," To demonstrate the difference between good and poor QSAA, the histograms of the 30 $\log g_{\rm e},T_{\rm e}$ values are plotted in the upper panels of Fig."886 2 for 4;=0.5.0.98 of SU Dra and BD |67 τος.," \ref{fig2} for $\varphi=0.5,0.98$ of SU Dra and BD +67 708."887 Phev show normal distributions with small scatter for the non-variable BD. |67 TOS and SU Dra at y=0.5., They show normal distributions with small scatter for the non-variable BD +67 708 and SU Dra at $\varphi=0.5$.888 At =0.98 the distribution is almost uniform., At $\varphi=0.98$ the distribution is almost uniform.889 At the next phase point 4=1 merely 14 intersections of US(loggoCh.Clos.MJ.ΕΣντο Le. 14. pairs of loggo.T; were found instead of 30 pairs.," At the next phase point $\varphi=1$ merely 14 intersections of $\{T_{\rm e}^{(i)}(\log g_{\rm e},{\rm CI}_1,{\rm CI}_2,[M],E(B-V))\}_890{i=1,2}$, i.e. 14 pairs of $\log g_{\rm e},T_{\rm e}$ were found instead of 30 pairs."891 Phus. at ος21 the observed. colours diller significantly from those of any static model of Ixurucz(1997).," Thus, at $\varphi \approx 1$ the observed colours differ significantly from those of any static model of \citet{kuru1}."892. Llowever. it is interesting to note that the small errors Aloeg.=0.03. AT.=191 indicate a phase island at y=0.93 with observed and static mocel colours in agreement for some 0.017zc10 minutes.," However, it is interesting to note that the small errors $\Delta\log g_{\rm e}=0.03$, $\Delta T_{\rm e}=12\mbox{K}$ indicate a phase island at $\varphi=0.93$ with observed and static model colours in agreement for some $0.01P\approx 10$ minutes."893 This happens to be the phase of the hump on the light curve when the inward and outward motions encounter and. produce a shock (Smith1995).., This happens to be the phase of the hump on the light curve when the inward and outward motions encounter and produce a shock \citep{smit1}.894 In the interval 0.92«45<1.05 the atmosphere is in a state of maximal compression by the shock coming from the sub-photospheric lavers and {1 is nearly minimal., In the interval $0.92 < \varphi < 1.05$ the atmosphere is in a state of maximal compression by the shock coming from the sub-photospheric layers and $R$ is nearly minimal.895 ‘This is the risine branch andthe start of the descending branch in the light curve., This is the rising branch andthe start of the descending branch in the light curve.896 Therefore. the values of logg..ἐν. obtained from QSSA. if they can be found at all. must be considered as a first approximation only.," Therefore, the values of $\log g_{\rm e},T_{\rm e},\vartheta$ obtained from QSSA, if they can be found at all, must be considered as a first approximation only."897 This is rellected in large Aloe q«. AZ. except for qz0.93.," This is reflected in large $\Delta \log g_{\rm e}$ , $\Delta T_{\rm e}$ except for $\varphi \approx 0.93$."898 To get raf) ete for (LOD). VO and Ώου.1). were differentiated by midpoint formulae.," To get $v(r,t)$ etc for \ref{107a}) ), $\vartheta(t)$ and $h_0(R,t)$ were differentiated by midpoint formulae."899 Fig., Fig.900 3aa is a plot of the functions ία.E). for the phases f= P. ge=0.15-0.35.0.5.0.55 with ai’!(4.40)=0.," \ref{fig3}a a is a plot of the functions ${\cal M}(d,t)$ for the phases $t=\varphi P$ , $\varphi=0.15\mbox{-}0.35,0.5,0.55$ with $a^{\rm (dyn)}(R,t)=0$."901 The average and standard vror of M.d are given in Table 2. from the pairs 25 and yo=0.15.0.3.0.35.0.55). (μι=0.35 and } in (10)) as our best. values denoted by —," The average and standard error of ${\cal M},d$ are given in Table \ref{tab2} from the pairs $(\varphi_1=0.25$ and $\varphi_2=0.15,0.3,0.35,0.55)$, $(\varphi_1=0.35$ and $\varphi_2=0.55)$ in \ref{107a}) ) as our best values denoted by $[\ast]$."902 yo=0.35 Condition Lis moderately violated but Condition LL is satisfied and a’(2.1D)gAR0x0.07. therefore. this phase was included in getting dM! of αι," At $\varphi=0.35$ Condition I is moderately violated but Condition II is satisfied and $a^{\rm (dyn)}(R,t)/g_{\rm903 s}(R,t)\approx -0.07$, therefore, this phase was included in getting $d,{\cal M}$ of $[\ast]$."904 Condition I is satisfied at 42=0.2.0.5. however. these phases had to be excluded from the mass and distance determination because of the large aU--=02.05)5.8.3.3ms7 and q=O52.38. respectively.," Condition I is satisfied at $\varphi=0.2,0.5$, however, these phases had to be excluded from the mass and distance determination because of the large $a^{\rm905 (dyn)}(R,\varphi=0.2,0.5)=5.8,-3.3\:\mbox{ms}^{-2}$ and $q=-0.52,38$, respectively."906 Using our best solution for .Vf ancld. the radius variation. velocities. and the components of acceleration were computed in physical units and are plotted inFig. 3bb-," Using our best solution for ${\cal M}$ and$d$, the radius variation, velocities, and the components of acceleration were computed in physical units and are plotted inFig.\ref{fig3}b b-d."907 Velocities ancl accelerations are plotted. only for. the phases of more or less good QSAA (ως 0.15-0.3.0.5.0.55). including the slightly shocked phases y= 0.35-0.45.," Velocities and accelerations are plotted only for the phases of more or less good QSAA $\varphi=0.15\mbox{-}0.3,0.5,0.55$ ), including the slightly shocked phases $\varphi=0.35\mbox{-}0.45$ ."908 At the phases «e=O.15.0.25.0.3.0.35.0.55 aPRu)=TOS.16.65.93ems7 and Condition Lis satisfied. [amenrp)(11D]< 0.13.Le. aPORus) is small incomparison with the other acceleration. ternis in (3).," At the phases $\varphi=0.15,0.25,0.3,0.35,0.55$ $a^{\rm (dyn)}(R,\varphi)=-79,8,-16,65,93\:\mbox{cms}^{-2}$ and Condition I is satisfied, $\vert a^{\rm (dyn)}(R,t)/g_{\rm s}(R,t)\vert < 0.13$ ,i.e. $a^{\rm (dyn)}(R,\varphi)$ is small incomparison with the other acceleration terms in \ref{1.100}) )."909 |g)&O.L is expected. from Alogg.= 0.04. δι=0.39. 0.02. 0.06. 0.10. 0.13. were. found.," $\vert q\vert\approx 0.1$ is expected from $\Delta \log g_{\rm e}=0.04$ , $q(R,\varphi)=0.39,$ $0.02,$ $0.06,$ $-0.10,$ $-0.13$ were found."910 Thus. Condition. Lb is indeed: satisfied.," Thus, Condition II is indeed satisfied."911Phe outlier value 0.39 is produced. by cancellation of OrfOl=—30.5mis ,The outlier value $0.39$ is produced by cancellation of $\partial v/\partial t\approx -30.5\:\mbox{ms}^{-2}$ 912thought it would be useful to check whether we could detect auv other edge. or check how strict the upper liuits are.,"thought it would be useful to check whether we could detect any other edge, or check how strict the upper limits are."913 We take NCC) 1051 as basis for comparison., We take NGC 4051 as basis for comparison.914 In NGC 1051. the absorber is rather Mehly ionized. with OVIII being the strongest edee witιτ LlicUl then NeX with r=0.8d0.1.," In NGC 4051, the absorber is rather highly ionized, with OVIII being the strongest edge with $\tau$ = $1.1 \pm 0.4$, then NeX with $\tau$ $0.8 \pm 0.4$."915 The OVII «dee is weaker with 720.35: I[xoniossa Fink. 1997.," The OVII edge is weaker with $\tau$ =0.35; Komossa Fink, 1997."916 Usiis the ratio Toy) TNX as typical for the relative clepths in a lighly ionized absorber. we can conclude that our non-doetection of edges other than the oue near 1.1 seV in PGIIOL]226 is still consistent with a standard wiwin absorber i.e. our upper luüts on 7 are not strict οrough to rule out a warm absorber m relativistic outflow.," Using the ratio $\tau$$_{OVIII}$ $\tau$$_{NeX}$ as typical for the relative depths in a highly ionized absorber, we can conclude that our non-detection of edges other than the one near 1.1 keV in PG1404+226 is still consistent with a standard warm absorber i.e. our upper limits on $\tau$ are not strict enough to rule out a warm absorber in relativistic outflow."917 Foursvectra were taken in February 1996 with IIST/FOS aud gcrafiues C130. CIO90II. (27) and CiLOO with iuteeration times 2300. 530. 120 aud 120 seconds respecively.," Four spectra were taken in February 1996 with HST/FOS and gratings G130, G190H, G270 and G400 with integration times 2300, 530, 120 and 120 seconds respectively."918 The total observed wavecheth range covered is 1087 - 1773 with a nominal resolution of 1300 (Fig.6)., The total observed wavelength range covered is 1087 - 4773 with a nominal resolution of 1300 (Fig.6).919 The spectra were taken through the 0.86 arcsec diameter aperture., The spectra were taken through the 0.86 arcsec diameter aperture.920 Standard reduction procedires were performed., Standard reduction procedures were performed.921 The wavelength scale of the spectrum taken with G19MI was shifted by |1.5A. to be consistent with the ot101) spectra.," The wavelength scale of the spectrum taken with G190H was shifted by +1.5, to be consistent with the other spectra."922 PC 11YL|226 was oserved with IUE ou Julv ] 2 davs after the ROSAT oervatious.," PG 1404+226 was observed with IUE on July 1994, 2 days after the ROSAT observations."923" The spectrum. SWP 51119. has a total integration time of 315r, ii (accumulated in 12 parts). and was taken through the large aperture aud at low dispersion."," The spectrum, SWP 51419, has a total integration time of 315 minutes (accumulated in 12 parts), and was taken through the large aperture and at low dispersion."924 Compared with the UST spectra taken LS mouthts later the coutimuui fiux iu July 1991 is zm 1.3 tinies brighter iu the common observed waveleugth range 1265 to but the Ίνα line kept the sale intensity., Compared with the HST spectra taken 18 monthts later the continuum flux in July 1994 is $\approx$ 1.3 times brighter in the common observed wavelength range 1265 to but the $\alpha$ line kept the same intensity.925 The modest S/N aud spectral POCsoll of the IUE spectra would xevent the detection of the absorption lines seen iu the IST spectra., The modest S/N and spectral resolution of the IUE spectrum would prevent the detection of the absorption lines seen in the HST spectra.926 No change iu the cuuissiou/absorption profile of Ίσα (Fie.7) cau be detected by comparing the IUE spectirmm aud the C130 spectrin rehbinued at2A., No change in the emission/absorption profile of $\alpha$ (Fig.7) can be detected by comparing the IUE spectrum and the FOS-G130 spectrum rebinned at.927 Table 3 lists the emission line imtensities (Ilj = 50 lau 1 Ἐν ου = 0. distance = 617 Mpce).," Table 3 lists the emission line intensities $_{0}$ = 50 km $^{-1}$ $^{-1}$, $_{0}$ = 0, distance = 617 Mpc)."928 We note the presence of some weak enission lues: (1) an unideutified line at (rest waveleneth LOTOA) noticed in a few other quasar spectra (Laor et al 1995: Taman et al 1997) (2) a line at 1290 wwhich we ideutifiv with 1175.7., We note the presence of some weak emission lines: (1) an unidentified line at (rest wavelength ) noticed in a few other quasar spectra (Laor et al 1995; Hamann et al 1997) (2) a line at 1290 which we identifiy with $^*$ 1175.7.929 This line is seen in UST spectra of Zwl and Laor et al (1997) sugeest that it is produced by resonance scattering of contimmuun photons by ions. a niechauis which requires large velocity eradicuts ( 1000 kins 1) within emitting cloud of the BLR.," This line is seen in HST spectra of IZw1 and Laor et al (1997) suggest that it is produced by resonance scattering of continuum photons by $^*$ ions, a mechanism which requires large velocity gradients $\approx$ 1000 km $^{-1}$ ) within emitting cloud of the BLR."930 We identify two absorption systems in the IIST. spectra which. in the source frame. are separated by ~1920 kin +.," We identify two absorption systems in the HST spectra which, in the source frame, are separated by $\sim 1920$ km $^{-1}$."931 Tn the blue svstem the absorption lines appear in the blue Πακ of the Ένα and CIV. emission lines at 800 luus ! from the peak., In the `blue' system the absorption lines appear in the blue flank of the $\alpha$ and CIV emission lines at 800 km $^{-1}$ from the peak.932 Iu the που system the absorption lues appear on the red flaik of the emission lines at 1100 kan + from the peak., In the `red' system the absorption lines appear on the red flank of the emission lines at 1100 km $^{-1}$ from the peak.933 Ilt is known that in radio quiet. ACN/QOuasars. the lieh ionization lines such as CTV are blucshitted with resect to the systemic velocity oa few hundred to a few thousand kin | (e.g. van Cronimgenun 1987. Corbin 1995. Sulentic et al.," It is known that in radio quiet AGN/Quasars, the high ionization lines such as CIV are blueshifted with respect to the systemic velocity by a few hundred to a few thousand km $^{-1}$ (e.g. van Groningen 1987, Corbin 1995, Sulentic et al."934 1995) this dueshiff beime generally iuerpreted as evidence for a wind outflowing from the facc' of the accretion disk turned owards us;, 1995) this blueshift being generally interpreted as evidence for a wind outflowing from the face of the accretion disk turned towards us.935 On this basis. we:weue that the red absorption," On this basis, we argue that the red absorption"936stars of varying mass and age.,stars of varying mass and age.937" With an uncertainty in [Fe/H] of +0.1, the masses of these stars could be determined to within 4 to7%."," With an uncertainty in $[\rm{Fe}/\rm{H}]$ of $\pm0.1$, the masses of these stars could be determined to within 4 to."938. Age can generally be determined to within 1 Gyr., Age can generally be determined to within 1 Gyr.939 The precision of the age is heavily influenced by theuncertainty in ὄνρο and the position of the star in the C-D diagram., The precision of the age is heavily influenced by theuncertainty in $\delta\nu_{02}$ and the position of the star in the C-D diagram.940 Many of these stars are being observed for an extended period of time withKepler and the precision of these measurements will undoubtedly improve., Many of these stars are being observed for an extended period of time with and the precision of these measurements will undoubtedly improve.941" With supporting ground-based spectroscopic observations and detailed modeling involving the individual frequencies, the fundamental properties of these stars will become well-determined and the data should provide significant tests of stellar models."," With supporting ground-based spectroscopic observations and detailed modeling involving the individual frequencies, the fundamental properties of these stars will become well-determined and the data should provide significant tests of stellar models."942 Figure 4 shows the observational e diagram., Figure \ref{fig2} shows the observational $\epsilon$ diagram.943" Observations cannot determine the radial order n directly, and so it is possible for ε to be uncertain by +1, particularly in the subgiants (Av in the range 20-80 "," Observations cannot determine the radial order $n$ directly, and so it is possible for $\epsilon$ to be uncertain by $\pm1$, particularly in the subgiants $\Delta\nu$ in the range 20–80 $\mu$ Hz)."944"For these, we have taken e to be in the range 0.7-1.7,wHz). but note there is some ambiguity for stars near the extremes of this range."," For these, we have taken $\epsilon$ to be in the range 0.7–1.7, but note there is some ambiguity for stars near the extremes of this range."945 The measurement of e is complicated by its close relationship to Av: a small change in Av can induce a large change in e., The measurement of $\epsilon$ is complicated by its close relationship to $\Delta\nu$: a small change in $\Delta\nu$ can induce a large change in $\epsilon$.946" We have also measured e using an alternative method: the variation of the large separation with frequency was measured as described by ?,, before the radial modes were globally fit and e derived, taking into account the mean curvature."," We have also measured $\epsilon$ using an alternative method: the variation of the large separation with frequency was measured as described by \citet{Mosser10}, before the radial modes were globally fit and $\epsilon$ derived, taking into account the mean curvature."947" À comparison of the values obtained by the two methods showed good agreement, although small systematic offsets exist, typically about 0.1."," A comparison of the values obtained by the two methods showed good agreement, although small systematic offsets exist, typically about 0.1."948 This offset is probably due to the combined effects of curvature (departure from equation (1))) and the slightly different range of frequencies over which Av and e were measured., This offset is probably due to the combined effects of curvature (departure from equation \ref{asymp}) )) and the slightly different range of frequencies over which $\Delta\nu$ and $\epsilon$ were measured.949 A single method must be applied to both models and data used to ensure consistency., A single method must be applied to both models and data used to ensure consistency.950 In this Letter we have used the method outlined in Section ??.., In this Letter we have used the method outlined in Section \ref{sec3}.951 The observed stars in Figure 4 are offset to the right of the models., The observed stars in Figure \ref{fig2} are offset to the right of the models.952" This offset is well-known from helioseismology, in which there is a discrepancy between the observed and computed oscillation frequencies of the Sun arising from improper modeling of the near-surface layers (??).."," This offset is well-known from helioseismology, in which there is a discrepancy between the observed and computed oscillation frequencies of the Sun arising from improper modeling of the near-surface layers \citep{Dziembowski88,C-D96}. ."953" A rigorous comparison of the observations with models requires that the offset be taken into account, either by a proper modeling of near-surface"," A rigorous comparison of the observations with models requires that the offset be taken into account, either by a proper modeling of near-surface"954sample shown in Fig.,sample shown in Fig.955 5 where the cross-correlation function peaks at an angular distance of 1. bubble radius with a significance of 9m., \ref{fig:crosscol} where the cross-correlation function peaks at an angular distance of 1 bubble radius with a significance of $\sigma$.956 In the immediate environment of a bubble the highest probability location to find an IMS YSO is projected against the rim of the bubble., In the immediate environment of a bubble the highest probability location to find an RMS YSO is projected against the rim of the bubble.957 Aloreover. it ds clear from inspecting Fig.," Moreover, it is clear from inspecting Fig."958 2. that the surface density of YSOs is not only enhanced at an angular olfset of1 bubble radius. but that it is enhanced over the entire. angular scale. of the bubbles out to an angular οκο of 2 bubble radii.," \ref{fig:surfdens} that the surface density of YSOs is not only enhanced at an angular offset of 1 bubble radius, but that it is enhanced over the entire angular scale of the bubbles out to an angular offset of 2 bubble radii."959" We can see this by comparing the mean surface clensity of YSOs ""μαρια 2 bubble radii ancl ""outside"" 2 bubble radii.", We can see this by comparing the mean surface density of YSOs “inside” 2 bubble radii and “outside” 2 bubble radii.960 The mean surface density of YSOs within an angular olfset of 2 bubble radii is S.O+1.7 YSOs/unit area compared a value of 3.20.2 YSOs/unit area at an angular ollset of 2 bubble radii or greater., The mean surface density of YSOs within an angular offset of 2 bubble radii is $\pm$ 1.7 YSOs/unit area compared a value of $\pm$ 0.2 YSOs/unit area at an angular offset of 2 bubble radii or greater.961 jX two sample unequal variance (heteroscedastic) t test of these two subsamples returns a probability. of only that these two subsamples are drawn from populations with the same mean., A two sample unequal variance (heteroscedastic) t test of these two subsamples returns a probability of only that these two subsamples are drawn from populations with the same mean.962 Hence we have demonstrated. that there is a statistically significant overdensity of massive YSOs associated with the bubbles compared to the hackground. with an enhanced. probability of finding these YSOs projected. against the rim. of the xibbles.," Hence we have demonstrated that there is a statistically significant overdensity of massive YSOs associated with the bubbles compared to the background, with an enhanced probability of finding these YSOs projected against the rim of the bubbles."963 What do these results imply?, What do these results imply?964 Firstly. there is a greater concentration of massive star formation towards the bubbles han in the wider environment.," Firstly, there is a greater concentration of massive star formation towards the bubbles than in the wider environment."965 This result is confirmed hy he surface density of MMD 6.7 Cillz masers (see Fig. 4)).," This result is confirmed by the surface density of MMB 6.7 GHz masers (see Fig. \ref{fig:mmb_surfdens}) ),"966 which trace a YSO population independently of mid-infrared emission., which trace a YSO population independently of mid-infrared emission.967 A greater concentration of star formation towards he bubbles implies that the bubbles are either cllicicnt at ooducing YSOs. or that they are found in regions of high YSO surface density.," A greater concentration of star formation towards the bubbles implies that the bubbles are either efficient at producing YSOs, or that they are found in regions of high YSO surface density."968 This is the classic chicken ancl egg scenario applied to massive star formation: do the bubbles »ecede the high surface density of YSOs. or does the high surface density of YSOs precede (or occur. sipultaneously with) the formation of the bubbles?," This is the classic chicken and egg scenario applied to massive star formation: do the bubbles precede the high surface density of YSOs, or does the high surface density of YSOs precede (or occur simultaneously with) the formation of the bubbles?"969 3elore considering this question more fully. we must ing in the second of our results that there is an enhanced oobabilitv of finding YSOs projected. against the rim of he bubbles (ie. at an angular olfset of 1: bubble radius)," Before considering this question more fully, we must bring in the second of our results – that there is an enhanced probability of finding YSOs projected against the rim of the bubbles (i.e. at an angular offset of 1 bubble radius)."970 Dv inspecting the autocorrelation of the IMS. YSOs we showed in Section 2.3 that this elfect is not likely to be due o intrinsic clustering within the RAIS sample on. similar angular scales to the bubble radii., By inspecting the autocorrelation of the RMS YSOs we showed in Section \ref{sect:autocol} that this effect is not likely to be due to intrinsic clustering within the RMS sample on similar angular scales to the bubble radii.971 Phe ancillary question raisecl by this result is: why are the YSOs more likely to be ound projected against the rim of the bubbles., The ancillary question raised by this result is: why are the YSOs more likely to be found projected against the rim of the bubbles.972 ie. what is special about the bubble rims?, i.e. what is special about the bubble rims?973 The bubble rims are traced by sym PATE emission which originates from the photon-clominatec region between the ionisation front being driven out by the LIL region. within the bubble and the surrounding neutral medium., The bubble rims are traced by $\mu$ m PAH emission which originates from the photon-dominated region between the ionisation front being driven out by the HII region within the bubble and the surrounding neutral medium.974 The rim of the bubbles thus shows the interface between LIL region and surrounding neutral gas., The rim of the bubbles thus shows the interface between HII region and surrounding neutral gas.975 For a spherical bubble morphology one would expect the column density of gas to be greater at the bubble rims due to the greater path length through the neutral material towards the rims., For a spherical bubble morphology one would expect the column density of gas to be greater at the bubble rims due to the greater path length through the neutral material towards the rims.976 So at first elance. the higher surface density of YSOs projected against the bubble rims may simply reflect. the higher column density at the rim of the bubbles. i.c. the YSOs trace molecular column density.," So at first glance, the higher surface density of YSOs projected against the bubble rims may simply reflect the higher column density at the rim of the bubbles, i.e. the YSOs trace molecular column density."977 However. while the sample of bubbles that have been observed at. relatively high angular resolution in molecular ines (27) do show a peaked molecular column density profile at à normalised) bubble radius of 1. the column cdensity alls olf much less sharply than the YSO surface. density.," However, while the sample of bubbles that have been observed at relatively high angular resolution in molecular lines \citep{beaumont2010} do show a peaked molecular column density profile at a normalised bubble radius of 1, the column density falls off much less sharply than the YSO surface density."978 Inspection of Figure 2 from ? shows that at a normalised xibble radius of 1.5 the CO intensity can be roughly half of hat at a normalised radius of L., Inspection of Figure 2 from \citet{beaumont2010} shows that at a normalised bubble radius of 1.5 the CO intensity can be roughly half of that at a normalised radius of 1.979 Phis suggests that the YSOs may not trace the column density distribution. although much closer scrutiny of the bubbles in a non-optically thick racer is required to confirm this hypothesis.," This suggests that the YSOs may not trace the column density distribution, although much closer scrutiny of the bubbles in a non-optically thick tracer is required to confirm this hypothesis."980 Moreover. the CO contrast between the centre of the bubbles and their rims is often extreme (2). whereas the YSO surface density within an angular olfset of 2 bubble radii is everywhere higher than the background level.," Moreover, the CO contrast between the centre of the bubbles and their rims is often extreme \citep{beaumont2010} whereas the YSO surface density within an angular offset of 2 bubble radii is everywhere higher than the background level."981 Vhus we cannot confidentIy. sav that the YSO surface density traces the gas column density around the bubbles., Thus we cannot confidently say that the YSO surface density traces the gas column density around the bubbles.982 The YSO surface density is strongly peaked at an olfset of 1 bubble radius and decreases sharply bevond this value., The YSO surface density is strongly peaked at an offset of 1 bubble radius and decreases sharply beyond this value.983 )evond an angular ollset of 2 bubble radii the surface density of. YSOs is essentially undistinguishable from the background. level., Beyond an angular offset of 2 bubble radii the surface density of YSOs is essentially undistinguishable from the background level.984 The angular cross-correlation function shows a similar steep drop —-| bevond an angular distance of 2 bubble radii the bubbles and RAIS YSOs are essentiallyuncorrelated., The angular cross-correlation function shows a similar steep drop — beyond an angular distance of 2 bubble radii the bubbles and RMS YSOs are essentially.985.. Vhe implication of this is that whatever causes the rise in YSO surface density is closely related to the rim of the bubbles., The implication of this is that whatever causes the rise in YSO surface density is closely related to the rim of the bubbles.986 The bubble radius is a dynamic value and expected to increase over time as stellar winds or radiation pressure causes the bubbles to expand., The bubble radius is a dynamic value and expected to increase over time as stellar winds or radiation pressure causes the bubbles to expand.987 Combined with this is the fact that the massive YSOs and. UC LIL regions identified by the RAIS survey typically tend to have lifetimes around a few 103 toa few 107 vears (2). and so should trace very recent star formation., Combined with this is the fact that the massive YSOs and UC HII regions identified by the RMS survey typically tend to have lifetimes around a few $^{4}$ to a few $10^{5}$ years \citep{mottram2011b} and so should trace very recent star formation.988 The sum of these pieces. of information leads us to conclude that it is likely that the bubbles predate the YSOs., The sum of these pieces of information leads us to conclude that it is likely that the bubbles predate the YSOs.989 Ifthe bubbles formed in an environment with a high surface density of YSOs (e.g.intheturbulenthighlyfragmentecinitialconditionssuggestedby 2).. then the distribution should not peak at the rim of the bubble as the bubble radius is time-dependent.," If the bubbles formed in an environment with a high surface density of YSOs \citep[e.g.~in the turbulent highly fragmented initial conditions suggested by][]{dale2011}, then the distribution should not peak at the rim of the bubble as the bubble radius is time-dependent."990 Similar arguments have been usec by ο for YSOs detected a he edges of shells in. Carina., Similar arguments have been used by \citet{preibisch2011} for YSOs detected at the edges of shells in Carina.991 Also in this case the extent of the enhanced YSO surface density should. also not be related to the current racius of the bubble why are bubbles found in regions of enhance YSO surface density occupying twice their angular radius?, Also in this case the extent of the enhanced YSO surface density should also not be related to the current radius of the bubble – why are bubbles found in regions of enhanced YSO surface density occupying twice their angular radius?992 Finally. the relative timescales of the massive YSOs and those required for the expansion of the bubbles imply hat the YSOs formed the bubbles.," Finally, the relative timescales of the massive YSOs and those required for the expansion of the bubbles imply that the YSOs formed the bubbles."993 We thus conclude hat a significant [fraction of the YSOs seen against the rim of the bubbles were likely triggered by the expansion of the rubble., We thus conclude that a significant fraction of the YSOs seen against the rim of the bubbles were likely triggered by the expansion of the bubble.994 A greater understanding of the clynamical timescales or the expansion of the bubbles and also the molecular environment of the bubbles are required. to confirm this ivpothesis., A greater understanding of the dynamical timescales for the expansion of the bubbles and also the molecular environment of the bubbles are required to confirm this hypothesis.995 Pinpointing the YSO formation to have occurred alter the bubble was formed is crucial to cisentaneling cause and effect in the star formation surrounding the bubbles., Pinpointing the YSO formation to have occurred after the bubble was formed is crucial to disentangling cause and effect in the star formation surrounding the bubbles.996 Currently. only a few bubbles have had their. dynamical lifetimes estimated and more studies similar to those of ? are required over a larger sample of bubbles.," Currently, only a few bubbles have had their dynamical lifetimes estimated and more studies similar to those of \citet{watson2009} are required over a larger sample of bubbles."997 Comparing the YSO clistribution to the gas distribution is also crucial to investigate dillerences in the population of YSOs at the rims of bubbles. for example to determine whether," Comparing the YSO distribution to the gas distribution is also crucial to investigate differences in the population of YSOs at the rims of bubbles, for example to determine whether"998CLUSTTERS OF GALAXIES CBenoist! Müunchen. Germany Copennhagen. Demnark The recent. discovery. of apparently massive ancl relaxed: clusters of galaxies at redshifts 220.5 ollers a unique opportunity to study the evolution of gravitationally bound svstems over an extended. look-back time.,"TERS OF GALAXIES }$, C. $^{1}$ unchen, Germany } nhagen, Denmark } The recent discovery of apparently massive and relaxed clusters of galaxies at redshifts $z\gsim0.5$ offers a unique opportunity to study the evolution of gravitationally bound systems over an extended look-back time."999 Moreover. if such systems are proven to be massive. especially those at + their sheer existence can impose stringent constraints on viable cosmological 20.8.modelsS].," Moreover, if such systems are proven to be massive, especially those at $z\gsim0.8$, their sheer existence can impose stringent constraints on viable cosmological models."1000 In addition. a large sample of confirmed clusters. spanning a broad redshift range. is of great interest for constraining moclels of formation and evolution of earlv-tv galaxies and scale structure. and for the selection of targetsin dillerent peredshift: intervals for," In addition, a large sample of confirmed clusters, spanning a broad redshift range, is of great interest for constraining models of formation and evolution of early-type galaxies and large-scale structure, and for the selection of targetsin different redshift intervals for"1001 (Rottecringetal.2003).. (Lonsdaleetal.2009) 2009).. (Oosterloo (Boothetal.2009):; (Jolustouctal.2008).," \citep{rbf03}, \citep{lcm+09} \citep{ecc+09}. \citep{Welch09,Dewdney09,Jonas09,ovc09}. \citep{ovc09}; \citep{bbjf09}; \citep{jtb+08}."1002. We divide radio transieuts iuto four categories based ou two attributes., We divide radio transients into four categories based on two attributes.1003 The first is the duration of the sje phenomenon (shorter than or greater than a few seconds)., The first is the duration of the basic phenomenon (shorter than or greater than a few seconds).1004 The secoud is their location (within the Calaxy or extra-galactic)., The second is their location (within the Galaxy or extra-galactic).1005 Roughly speaking the duration maS o coherent versus incohercut cliission aud the location o repeated versus catacbesinice eveuts., Roughly speaking the duration maps to coherent versus incoherent emission and the location to repeated versus cataclysmic events.1006 Pulsars aud related phenomenon (eiut pulses. nulli18o oulsus. erratic pulsars. roating radio transieuts. aid uagnetars) are the dominant category of short duration radio trausieuts at meter and centimneter wavoleusths.," Pulsars and related phenomenon (giant pulses, nulling pulsars, erratic pulsars, rotating radio transients, and magnetars) are the dominant category of short duration radio transients at meter and centimeter wavelengths."1007 There are no secure examples of short duration radio rausicuts that are located bevoud the local Group., There are no secure examples of short duration radio transients that are located beyond the local Group.1008 Flare stars and associated phenomena are prine examples of oue duration radio transicuts of Galactic origin., Flare stars and associated phenomena are prime examples of long duration radio transients of Galactic origin.1009 The focus of this paper is long duration trausicuts of extra-galactic origin., The focus of this paper is long duration transients of extra-galactic origin.1010 Known examples iu this group are supernovae (Weileretal.2010) aud ginuuarav burst afterelows (Gehrelsetal.2009).," Known examples in this group are supernovae \citep{wps+10}1011 and gamma-ray burst afterglows \citep{grf09}."1012. Iu both cases. the radio eniüssion arises as the fast moving debris mteracts with the circustellar matter.," In both cases, the radio emission arises as the fast moving debris interacts with the circumstellar matter."1013 Ii Table 1 we «παχο the areal density of racio-cuuitting supernovae (iucludiug the sub-classes) aud GRB afterglows., In Table \ref{tab:ListOfTrans} we summarize the areal density of radio-emitting supernovae (including the sub-classes) and GRB afterglows.1014" Note the areal density of ""live transicuts” (transicuts present at anv even instant of time) of both supernovae and CRB afterelows is less than 0.05 per square degree.", Note the areal density of “live transients” (transients present at any given instant of time) of both supernovae and GRB afterglows is less than 0.05 per square degree.1015 Iu 2007. Bower ct reported on the analysis of a single feld observed hereafter.every BOT]week as a part of the Very Large Axrav. (VLA) calibration progriun.," In 2007, Bower et \nocite{bsb+07} [hereafter, B07] reported on the analysis of a single field observed every week as a part of the Very Large Array (VLA) calibration program."1016 The observations were conducted at GGIIz aud CCTz and lasted 22 wears., The observations were conducted at GHz and GHz and lasted 22 years.1017 The 9114 epochs and the lLweekly cadence inakes this data set a most valuable set to probe the decimeter baud for loug duration trausieuts at the subanilliJausky level., The 944 epochs and the weekly cadence makes this data set a most valuable set to probe the decimeter band for long duration transients at the sub-milliJansky level.1018" These authors reported the discovery of eight transicuts found iu only one epoch (hereafter. ""iugle-epoch: duration. 20nuuimutes"," These authors reported the discovery of eight transients found in only one epoch (hereafter “single-epoch”; duration, minutes"1019where didt=OfOF|veNV ods the Lagrangian time derivative. £4 is the CR pressure. ος is the CR energy density. & is the CR cliffsion coefficient. 5; is the CR source terii due to the central AGN activity. pg. is the iron density. and all other variables have their usual meanings.,"where $d/dt \equiv \partial/\partial t+{\bf v} \cdot \nabla $ is the Lagrangian time derivative, $P_{\rm c}$ is the CR pressure, $e_{\rm1020 c}$ is the CR energy density, $\kappa$ is the CR diffusion coefficient, $\dot{S_{\rm c}}$ is the CR source term due to the central AGN activity, $\rho_{\rm Fe}$ is the iron density, and all other variables have their usual meanings."1021" Pressures aud euergy densities are related via P=(>De aud BP.=(τιle. where we assume ~=h/8 ands,=1/3."," Pressures and energy densities are related via $P=(\gamma-1)e$ and $P_{\rm c}=(\gamma_{\rm c}-1)e_{\rm c}$, where we assume $\gamma=5/3$ and $\gamma_{\rm c}=4/3$."1022 Equation (5)) describes the couservation of iron mass., Equation \ref{hydro5}) ) describes the conservation of iron mass.1023 Since we focus ou the effect of ACN outbursts on tle iron distribution and follow the cluster evolution for a timescale much shorter than eurichinoeut times (2.5 Cor: 7)). we ignore the iron source term.," Since we focus on the effect of AGN outbursts on the iron distribution and follow the cluster evolution for a timescale much shorter than enrichment times $\gtrsim 5$ Gyr; \citealt{bohringer04}) ), we ignore the iron source term."1024 The iron abuudance Z iu unitsof the solar value is proportional to PE/p., The iron abundance $Z$ in unitsof the solar value is proportional to $Z\propto \rho_{\rm Fe}/\rho$ .1025 Thus the iron density pg. in equation (5)) may be replaced by Zp., Thus the iron density $\rho_{\rm Fe}$ in equation \ref{hydro5}) ) may be replaced by $Z \rho$.1026 Siuce both the gas mass aud irou lass are conserved. the metallicity Z is also conserved: and therefore provides a tracer for the ICM eas. which is helpful in understanding how ACN outburstsaffect and mix the ICM.," Since both the gas mass and iron mass are conserved, the metallicity $Z$ is also conserved: and therefore provides a tracer for the ICM gas, which is helpful in understanding how AGN outburstsaffect and mix the ICM."1027 Iu the eas energv equation (3)). we inchide radiative cooling with a volume cooling rate nye(T.Z). where the cooling function A(7.Z) is adopted from ?.— and depends ou both gas temperature TZ and metallicity Z.," In the gas energy equation \ref{hydro3}) ), we include radiative cooling with a volume cooling rate $n_{\rm i}n_{\rm e}\Lambda(T,Z)$, where the cooling function $\Lambda(T,Z)$ is adopted from \citet{sd93}1028 and depends on both gas temperature $T$ and metallicity $Z$."1029 The jou umber deusitv 5»; is related to the proton nuuber deusifv ay via sj= Lloeg. aud thus the molecular weight is 4—0.61., The ion number density $n_{\rm i}$ is related to the proton number density $n_{\rm H}$ via $n_{\rm i}=1.1n_{\rm H}$ and thus the molecular weight is $\mu=0.61$.1030" The gas temperature is related to the gas pressure aud density via the ideal gas law: where Ay is Boltzmann's coustaut aud i, is the atomic lass unt.", The gas temperature is related to the gas pressure and density via the ideal gas law: where $k_{\rm B}$ is Boltzmann's constant and $m_{\mu}$ is the atomic mass unit.1031 Equatious (1)) (5)) ave solved im (r0:) evlindrica coordinates using a two-dimensional Eulerian code «λίαν to ZEUS 2D να iu particular. we have incorporated iuto the code a background eravitationa potential. CR diffusion. CR cucrey equation. aud mon equation of mass conservation.," Equations \ref{hydro1}) ) $-$ \ref{hydro5}) ) are solved in $(r, z)$ cylindrical coordinates using a two-dimensional Eulerian code similar to ZEUS 2D \citep{stone92}; ; in particular, we have incorporated into the code a background gravitational potential, CR diffusion, CR energy equation, and iron equation of mass conservation."1032 The computational eric consists of 100 equally spaced zones in both coordinates out to 100 kpe plus additional 100 logarithiuically-spacec zoues out to 1 Alpe., The computational grid consists of $100$ equally spaced zones in both coordinates out to $100$ kpc plus additional $100$ logarithmically-spaced zones out to $1$ Mpc.1033 For all the three fluids. we adopt reflective boundary conditions at the origin aud outflow boundary conditions at the outer boundary.," For all the three fluids, we adopt reflective boundary conditions at the origin and outflow boundary conditions at the outer boundary."1034 Our model aud methods are geucrallv applicable to all relaxed clusters. but for concreteuess. we adopt «λαο. paraimcters appropriate for the typical CC cluster Abell 1795. which has been well observed by both aud (2???) ," Our model and methods are generally applicable to all relaxed clusters, but for concreteness, we adopt simulation parameters appropriate for the typical CC cluster Abell 1795, which has been well observed by both and \citep{tamura01, ettori02, vikhlinin06}. ."1035"For initial profiles of AL795. we first build au analytic fit to the deprojected 3-dimnieunsional eas temperature profile derived. from observations. which acquired data out to ~1 Mpe covering our eutire coniputational eril: where the constant a is chosen to be a= 5. T,=2.5 keV is the observed central temperature of A1795.and Ty isthebest-fit temiperature profile of ?. which providesan excellent fit to Chandra data of AT795 from LO kpe to l Mpc: where"," For initial profiles of A1795, we first build an analytic fit to the deprojected $3$ -dimensional gas temperature profile derived from observations, which acquired data out to $\sim 1$ Mpc covering our entire computational grid: where the constant $\alpha$ is chosen to be $\alpha=5$ , $T_{\rm in}=2.5$ keV is the observed central temperature of A1795,and $T_{\rm V}$ is thebest-fit temperature profile of \citet{vikhlinin06} which providesan excellent fit to data of A1795 from $40$ kpc to $1$ Mpc: where"1036wave equation: where the operator PF is given by Here z is the distance to the magnetic axis of symmetry.,wave equation: where the operator $F$ is given by Here $x$ is the distance to the magnetic axis of symmetry.1037" Although in the presence of a mixed poloidal and toroidal field the equations still give rise to a continuous set of solutions, the calculations are significantly complicated as the continuum modes are affected by the toroidal component of the field, by gravity, and by compressibility."," Although in the presence of a mixed poloidal and toroidal field the equations still give rise to a continuous set of solutions, the calculations are significantly complicated as the continuum modes are affected by the toroidal component of the field, by gravity, and by compressibility."1038 For the sake of simplicity we will ignore toroidal fields in our dynamic simulations., For the sake of simplicity we will ignore toroidal fields in our dynamic simulations.1039" We will however, calculate the continuum frequencies for a mixed poloidal and toroidal field in the Appendix B. For determining the spectrum of the core continuum, the appropriate boundary conditions are €4(x=xc)0, where χε(Φ) marks the location of the crust-core interface."," We will however, calculate the continuum frequencies for a mixed poloidal and toroidal field in the Appendix B. For determining the spectrum of the core continuum, the appropriate boundary conditions are $\xi_{\phi}(\chi=\chi_c)=0$, where $\chi_c(\phi)$ marks the location of the crust-core interface."1040" With this boundary condition, Equation (26)) constitutes a Sturm-Liouville problem on each separate flux surface i»."," With this boundary condition, Equation \ref{poedts}) ) constitutes a Sturm-Liouville problem on each separate flux surface $\psi$."1041" Using the stellar structure model and magnetic field configuration from section 4.1, we can calculate the eigenfunctions and eigenfrequencies for each flux surface w."," Using the stellar structure model and magnetic field configuration from section 4.1, we can calculate the eigenfunctions and eigenfrequencies for each flux surface $\psi$."1042 The reflection symmetry of the stellar model and the magnetic field with respect to the equatorial plane assures that the eigenfunctions of Eq. (26)), The reflection symmetry of the stellar model and the magnetic field with respect to the equatorial plane assures that the eigenfunctions of Eq. \ref{poedts}) )1043 are either symmetric or anti-symmetric with respect to the equatorial plane., are either symmetric or anti-symmetric with respect to the equatorial plane.1044 We can therefore determine the eigenfunctions by integrating Eq. (26)), We can therefore determine the eigenfunctions by integrating Eq. \ref{poedts}) )1045 along the magnetic field lines from the equatorial plane x=0 to the crust-core interface x=xc(v).," along the magnetic field lines from the equatorial plane $\chi =10460$ to the crust-core interface $\chi = \chi_c \left( \psi \right)$."1047" Let us consider the odd modes here for which £5(0)=0, and solve Eq. (26))"," Let us consider the odd modes here for which $\xi_{\phi} \left( 0 \right) = 0$, and solve Eq. \ref{poedts}) )"1048" with the boundary condition £5(x.)=0 at the crust-core interface; for even modes, the boundary condition is d£;(0)/dx=0."," with the boundary condition $\xi_{\phi} \left(\chi_c \right) =10490$ at the crust-core interface; for even modes, the boundary condition is $d\xi_{\phi} \left( 0 \right)/d\chi=0$."1050 We find the eigenfunctions by means of a shooting method; using fourth order Runge-Kutta integration we integrate from x=0 to x=x..," We find the eigenfunctions by means of a shooting method; using fourth order Runge-Kutta integration we integrate from $\chi1051= 0$ to $\chi = \chi_c$."1052" The correct eigenvalues o,, and eigenfunctions £,(x) are found by changing the value of o until the boundary condition at En is satisfied.", The correct eigenvalues $\sigma_n$ and eigenfunctions $\xi_n \left( \chi \right)$ are found by changing the value of $\sigma$ until the boundary condition at $\xi_n$ is satisfied.1053 In this way we gradually increase the value of σ until the desired number of harmonics is obtained., In this way we gradually increase the value of $\sigma$ until the desired number of harmonics is obtained.1054 In figure 12 we show a typical resulting core-continuum.," In figure \ref{core_cont1}1055 we show a typical resulting core-continuum."1056" According to Sturm-Liouville theory the normalized eigenfunctions £, of Eq. (26))", According to Sturm-Liouville theory the normalized eigenfunctions $\xi_n$ of Eq. \ref{poedts}) )1057 form an orthonormal basis with respect to the following inner product:, form an orthonormal basis with respect to the following inner product:1058the conventional law Τ(λ)οςA7? (Draine1989).,the conventional law $\tau (\lambda) \propto \lambda^{-1.75}$ \citep{Dr89}.1059". In this assumption, the total flux density is written as follows, where rp and ηλαν are the flux density amplitudes of the starburst and AGN templates {235 and f29 normalized at 6 um, respectively."," In this assumption, the total flux density is written as follows, where $\eta_{\rm SB}$ and $\eta_{\rm AGN}$ are the flux density amplitudes of the starburst and AGN templates $f^{\rm SB}_{\nu}$ and $f^{\rm AGN}_{\nu}$ normalized at 6 $\mu$ m, respectively."1060" We can estimate the only two free parameters (the ratio of sp to ηλαν and το, which is the 6 pom optical depth to the AGN) by fitting the spectrum."," We can estimate the only two free parameters (the ratio of $\eta_{\rm SB}$ to $\eta_{\rm AGN}$ and $\tau_{6}$, which is the 6 $\mu$ m optical depth to the AGN) by fitting the spectrum."1061" From these, we can get the intrinsic AGN contribution to the 6 jum flux density, ag=nacn/(Nacn+risp)."," From these, we can get the intrinsic AGN contribution to the 6 $\mu$ m flux density, $\alpha_{6} = \eta_{\rm AGN} / (\eta_{\rm AGN} + \eta_{\rm SB})$."1062" In addition, we can also estimate the AGN contribution to the total infrared luminosity (this roughly corresponds to the bolometric luminosity, Loi, for ULIRGs), oo=nacn/(Nacnt+Knsp), where K=RACN/RS®, and RAGN and R?P are the ratios of the intrinsic flux at 6 wm to the total infrared flux of AGN and starbursts, respectively; in local ULIRGs, logRAGN= logRSP=1005, yielding K~35 (Nardini et al."," In addition, we can also estimate the AGN contribution to the total infrared luminosity (this roughly corresponds to the bolometric luminosity, $L_{\rm bol}$, for ULIRGs), $\alpha_{\rm bol} = \eta_{\rm AGN} / (\eta_{\rm AGN} + K \eta_{\rm SB}) $, where $K = R^{\rm AGN}/R^{\rm SB}$, and $R^{\rm AGN}$ and $R^{\rm SB}$ are the ratios of the intrinsic flux at 6 $\mu$ m to the total infrared flux of AGN and starbursts, respectively; in local ULIRGs, $\log R^{\rm AGN} = -0.36^{+0.06}_{-0.07}$ , $\log R^{\rm SB} = -1.91^{+0.02}_{-0.02}$, yielding $K \sim 35$ (Nardini et al."1063" in —0.36*905,prep.).", in prep.).1064" The model--1.91 has been successfully applied to a sample of local ULIRGs, obtaining two important results: 1) despite the complexity of the spectra, all the sources were successfully fitted, showing that the relative AGN/starburst contribution and the extinction of the AGN component are responsible for most of the observed variety; 2) the expected AGN/starburst contributions to the bolometric luminosity reproduce closely the total observed luminosity."," The model has been successfully applied to a sample of local ULIRGs, obtaining two important results: 1) despite the complexity of the spectra, all the sources were successfully fitted, showing that the relative AGN/starburst contribution and the extinction of the AGN component are responsible for most of the observed variety; 2) the expected AGN/starburst contributions to the bolometric luminosity reproduce closely the total observed luminosity."1065" This is an important test for our model: the bolometric contributions are obtained from the 6 jum spectral decomposition only (and from the average bolometric ratios, which have a fixed value for all objects), therefore the comparison between the predicted and observed total luminosities is an independent check of the results."," This is an important test for our model: the bolometric contributions are obtained from the 6 $\mu$ m spectral decomposition only (and from the average bolometric ratios, which have a fixed value for all objects), therefore the comparison between the predicted and observed total luminosities is an independent check of the results."1066 Here we apply the same model to high-redshift sources., Here we apply the same model to high-redshift sources.1067" In doing so, we assume that a) the intrinsic SED, and b) the bolometric ratios, are the same at low and high redshift."," In doing so, we assume that a) the intrinsic SED, and b) the bolometric ratios, are the same at low and high redshift."1068 We discuss the implications and the limits of these assumptions in 84., We discuss the implications and the limits of these assumptions in 4.1069 The stacked spectra for the two samples of submm- and 24 um-selected galaxies are shown in Fig., The stacked spectra for the two samples of submm- and 24 $\mu$ m-selected galaxies are shown in Fig.1070" 1, together with our deconvolution in the AGN and starburst components."," 1, together with our deconvolution in the AGN and starburst components."1071in the FGM-based BLFs (see the closed contours in Fig. 39).,in the FGM-based BLFs (see the closed contours in Fig. \ref{fig:gauss_lf}) ).1072 This structure is introduced by the Gaussian copula. and from the physical point of view. it might not be strongly desired.," This structure is introduced by the Gaussian copula, and from the physical point of view, it might not be strongly desired."1073 The FGM-based BLF has a more ideal shape., The FGM-based BLF has a more ideal shape.1074 Second. since the univariate LF shapes are different at FIR and FUV. the ridge of the BLF is not a straight line but clearly nonlinear.," Second, since the univariate LF shapes are different at FIR and FUV, the ridge of the BLF is not a straight line but clearly nonlinear."1075 This feature is more clearly visible in higher correlation cases in Figure 3.. but always exists for the whole range of p.," This feature is more clearly visible in higher correlation cases in Figure \ref{fig:gauss_lf}, but always exists for the whole range of $\rho$."1076 This trend is indeed found in the {ην diagram (Martinetal.2005)., This trend is indeed found in the $\lir$ $\luv$ diagram \citep{martin05}.1077. The underlying physies is that galaxies with high SFRs are more extinguished by dust (e.g.Buatetal.2007a.b).," The underlying physics is that galaxies with high SFRs are more extinguished by dust \citep[e.g.][]{buat07a,buat07b}."1078. Observational applications including this topic will be presented elsewhere (Takeuchi et 2010. in preparation).," Observational applications including this topic will be presented elsewhere (Takeuchi et 2010, in preparation)."1079 Since we have an explicit form of a BLF. we can discuss the flux selection effect formally.," Since we have an explicit form of a BLF, we can discuss the flux selection effect formally."1080 For simplicity. we consider the bivariate case ssample selected at two bands). but it will be straightforward to extend the formulation to a multiwavelength ease (or more generally. selected using any physical properties).," For simplicity, we consider the bivariate case sample selected at two bands), but it will be straightforward to extend the formulation to a multiwavelength case (or more generally, selected using any physical properties)."1081" The flux selection is described in terms of luminosity as putting a lower bound LP"" on a CL —L:2) plane.", The flux selection is described in terms of luminosity as putting a lower bound $\llim$ on a luminosity--luminosity $L_1$ $L_2$ ) plane.1082" The lower bound luminosity LE"" is defined by the flux (density) detection limit SI?"" as a function of redshift.", The lower bound luminosity $\llim$ is defined by the flux (density) detection limit $S^{\rm lim}$ as a function of redshift.1083 In most surveys. a certain wavelength band is chosen as the primary selection band. like B-band. Ks-band. 60jm-selected. ete.," In most surveys, a certain wavelength band is chosen as the primary selection band, like -band, s-band, $60\;\mu$ m-selected, etc."1084 The schematic description of a survey is presented in Figure 4.., The schematic description of a survey is presented in Figure \ref{fig:selection_effect}.1085" If we select a sample of objects (in our case galaxies) at band |. the objects with £L,<LY""(73 would not be included in the sample at a certain redshift z."," If we select a sample of objects (in our case galaxies) at band 1, the objects with $L_1 < \llim_1 (z)$ would not be included in the sample at a certain redshift $z$."1086" Then. the detected sources should have £1)=Li"" zÉj)and Lo=nu2)."," Then, the detected sources should have $L_1 > \llim_1 (z)$ and $L_2 > \llim_2 (z)$."1087" Hence. on the £,—L» plane. the 2-dim distribution of the detected objects is expressed as where © is a solid angle. and © is the Heaviside step function detined as The quantity S° is proportional to the surface number density of objects detected in both bands on the {ιο plane."," Hence, on the $L_1$ $L_2$ plane, the 2-dim distribution of the detected objects is expressed as where $\Omega$ is a solid angle, and $\Theta$ is the Heaviside step function defined as The quantity $\sdet$ is proportional to the surface number density of objects detected in both bands on the $L_1$ $L_2$ plane."1088 We start from a primary selection at band |. then we would have objects detected at band | but not detected at band 2.," We start from a primary selection at band 1, then we would have objects detected at band 1 but not detected at band 2."1089 In such a case we only have upper limits for these objects., In such a case we only have upper limits for these objects.1090" The 2-dim distribution of the upper limits at band 2 is similarly formulated as The superseript UL? stands for ""upper limit at band 27.", The 2-dim distribution of the upper limits at band 2 is similarly formulated as The superscript UL2 stands for “upper limit at band 2”.1091 In statistical terminology. the upper limit case. wwe know there is an object but we do only have the upper tor lower) limits of a certain quantity. is referred to as “censored”.," In statistical terminology, the upper limit case, we know there is an object but we do only have the upper (or lower) limits of a certain quantity, is referred to as “censored”."1092 Though we can detine the distribution XUL.Es) by Eq. (09. ," Though we can define the distribution $\sult (L_1, L_2)$ by Eq. \ref{eq:ul2}) ),"1093since the sample objects belonging to this category appear only as upper limits on the plot. a special statistical treatment. referred to as the survival analysis. is required to estimate »ial£2) from the data.," since the sample objects belonging to this category appear only as upper limits on the plot, a special statistical treatment, referred to as the survival analysis, is required to estimate $\sult (L_1, L_2)$ from the data."1094 Since we select objects at band |. we do not have upper limits at band |. because we do not know if there would be an object below the limit.," Since we select objects at band 1, we do not have upper limits at band 1, because we do not know if there would be an object below the limit."1095" This case is called ""truncated"" in statistics.", This case is called “truncated” in statistics.1096 If we select objects at band 2. we can formulate the 2-dim distribution of detected objects and upper limits exactly in the same way as the band | selected sample.," If we select objects at band 2, we can formulate the 2-dim distribution of detected objects and upper limits exactly in the same way as the band 1 selected sample."1097 For the objects detected at both bands. the 2-dim distribution is expressed by Eq. (38)).," For the objects detected at both bands, the 2-dim distribution is expressed by Eq. \ref{eq:detect}) )."1098 The objects detected at band 2 but not detected at band | is expressed as ↕↑⊲∖↖∁∁⋡∙⋯∣∏⋯∣∁∣∠↕⊥⊔⊔↿∖∶⊐⋡∙⋯↳∣∠⊳↕∙↘⊔⊔↿∖∶⊐⇂≯⇂⊾∁∁↥⊰∁∣⋝⊽↥⋂∁∣⊔↳∐⋂∙⋮↾↾∣↧∁∫∖⋡−∁⋯⊾⇂⊾∁∁⊓∩⋂⋅∁∖⊽∩∣⊔⊓⋯⊤∙∥⊾⋝⊽∁⇈⊲∁∁↾⋅∁↾∁⋅⋅∖↖," The objects detected at band 2 but not detected at band 1 is expressed as If we can model $\llim_1 (z)$ and $\llim_2 (z)$ precisely including the $K$ -correction, evolutionary effect, etc.,"1099∁∁⋡∙⋯⊔⊰∁↾∣↧∁∩∣↴⊰∁∏⊽∁↲∣↴↥∖∩∙∣∏⋡∙∣↾∁ luminosity distribution to estimate the correlation coefficient. or more generally the dependence structure of two luminosities through Eqs. (38))-(40)).," we can use the observed bivariate luminosity distribution to estimate the correlation coefficient, or more generally the dependence structure of two luminosities through Eqs. \ref{eq:detect}) \ref{eq:ul2}) )."1100 We can deal with these cases in a unified manner with techniques developed in survival analysis., We can deal with these cases in a unified manner with techniques developed in survival analysis.1101 We discuss this issue ina subsequent work (Takeuchi et 22010. in preparation).," We discuss this issue in a subsequent work (Takeuchi et 2010, in preparation)."1102 The star formation rate (SFR) is one of the most fundamental quantities to investigate the formation and evolution of galaxies., The star formation rate (SFR) is one of the most fundamental quantities to investigate the formation and evolution of galaxies.1103 The SFR is often estimated from the FUV flux of galaxies (or other related observables like Ha ete.), The SFR is often estimated from the FUV flux of galaxies (or other related observables like $\alpha$ etc.)1104" after ""correcting"" the dust extinction.", after “correcting” the dust extinction.1105 However.," However,"1106appropriately.,appropriately.1107 The number of stars per magnitude (7) and colour interval (7. .7) was integrated to the magnitude limit of this survey., The number of stars per magnitude $I$ ) and colour interval $I-J$ ) was integrated to the magnitude limit of this survey.1108 The derived M-dwarf isonumbers are compared to our data in Fig. 6.., The derived M-dwarf isonumbers are compared to our data in Fig. \ref{Figcontamination}.1109 The conclusion is that it is not likely that the proposed new Pleiades BDs are field M dwarfs., The conclusion is that it is not likely that the proposed new Pleiades BDs are field M dwarfs.1110 Eight new Pletades candidates have been identified. four of which are possible BDs.," Eight new Pleiades candidates have been identified, four of which are possible BDs."1111 Three of the four brightest new candidates have proper motions consistent with Pletades membership (Hambly. priv.," Three of the four brightest new candidates have proper motions consistent with Pleiades membership (Hambly, priv."1112 comm.)., comm.).1113 Two probable members (NPL22 32) stick out from the single-star sequence and are analyzed as binaries together with the spectroscopic binary PPL15 (NPL35)., Two probable members (NPL22 32) stick out from the single-star sequence and are analyzed as binaries together with the spectroscopic binary PPL15 (NPL35).1114 À number of faint very red objects were also found., A number of faint very red objects were also found.1115 Two of those were measured also in A and show colours similar to GDI65B and are possible field BDs., Two of those were measured also in $K$ and show colours similar to GD165B and are possible field BDs.1116 In Fig., In Fig.1117 7. this survey 1s compared to several other recent surveys., \ref{Figcomparison} this survey is compared to several other recent surveys.1118 Known nonmembers have been excluded., Known nonmembers have been excluded.1119 The dispersion of the Steele et al., The dispersion of the Steele et al.1120 (1993.1995) data can probably be explained by photometric uncertainty. since most of their I magnitudes are photographic.," \cite*{steele93,steele95} data can probably be explained by photometric uncertainty, since most of their $I$ magnitudes are photographic."1121 Note that the faint Pleiades sequence is slightly bluer than the Baraffe et al., Note that the faint Pleiades sequence is slightly bluer than the Baraffe et al.1122 (1998) model., \cite*{baraffe98} model.1123 Part of this may be due to incomplete line lists and not yet included dust formation in the models., Part of this may be due to incomplete line lists and not yet included dust formation in the models.1124 Note also that Mermilliod et al., Note also that Mermilliod et al.1125 (1997) found from Hippareos data that the Pleiades cluster is peculiar in the sense that its main sequence is ~0.1 mag fainter than other nearby clusters. such as the Hyades and Praesepe.," \cite*{mermilliod97} found from Hipparcos data that the Pleiades cluster is peculiar in the sense that its main sequence is $\sim0.4$ mag fainter than other nearby clusters, such as the Hyades and Praesepe."1126 The reason for this peculiarity is not known. and may also hide part of the model deviation.," The reason for this peculiarity is not known, and may also hide part of the model deviation."1127 The objects in Table 2 that are of special interest are individually discussed and compared to other papers below.N," The objects in Table \ref{Tabphotometry} that are of special interest are individually discussed and compared to other papers below.,"1128PL11. and have proper motions consistent with membership (Hambly. priv.," and have proper motions consistent with membership (Hambly, priv."1129 comm.).," comm.),"1130 although not present in HHJ., although not present in HHJ.1131 NPL22 is also a possible binary. best fitted by two components of equal brightness. Z4=Ij16.3NPL24. and have proper motions that are not consistent with membership. (," NPL22 is also a possible binary, best fitted by two components of equal brightness, $I_\mathrm{A}=I_\mathrm{B}=16.3$, and have proper motions that are not consistent with membership. ("1132ΡΡΙ12) and have uncertain proper motions.,PPL12) and have uncertain proper motions.1133 NPL30 has a radial velocity consistent with membership (Staufferetal.1994b)., NPL30 has a radial velocity consistent with membership \cite{stauffer94b}.1134. NPL32 is very close to a bright star. which due to blending makes the photographie proper motion uncertain.," NPL32 is very close to a bright star, which due to blending makes the photographic proper motion uncertain."1135 If NPL32 is à member. its position above the Pleiades sequence indicates an unresolved binary. best fitted by Iq =17.1. fy=15.0. (," If NPL32 is a member, its position above the Pleiades sequence indicates an unresolved binary, best fitted by $I_\mathrm{A}=17.4$ , $I_\mathrm{B}=18.0$. ("1136HHJ26) is. as also found by Steele et al. (1993).,"HHJ26) is, as also found by Steele et al. \cite*{steele93},"1137. an RE nonmember. (, an $RI$ nonmember. (1138PPLI5) has been measured by several authors recently (Stauffer et al.,PPL15) has been measured by several authors recently (Stauffer et al.1139 1994a: Basri et al., 1994a; Basri et al.1140" 1996: ZMR). and also found to be a spectroscopic binary (Basri&Martin1997).,From the primary component's possible loeit in our colour-magnitude diagrams. the secondary’s mass is ~0.03 M... consistent with ZMR."," 1996; ZMR), and also found to be a spectroscopic binary \cite{basri97}.From the primary component's possible locii in our colour-magnitude diagrams, the secondary's mass is $\sim0.03$ $_{\odot}$, consistent with ZMR."1141 The 7 magnitudes would be 7/4=17.96 Hm=21.3., The $I$ magnitudes would be $I_\mathrm{A}=17.96$ $I_\mathrm{B}=21.3$.1142 A heavier and brighter secondary would force the primary below the disk sequence., A heavier and brighter secondary would force the primary below the disk sequence.1143 are all below the BD limit., are all below the BD limit.1144 NPL37 shows a slight brightness enhancement at the edge of the stellar profile., NPL37 shows a slight brightness enhancement at the edge of the stellar profile.1145 It is not clear wether this is a background star or galaxy or if NPL37 itself i$ a compact galaxy. (, It is not clear wether this is a background star or galaxy or if NPL37 itself is a compact galaxy. (1146Teidel).,Teide1).1147 Our result is 7—19.26. 7=16.15.," Our result is $I=19.26$, $J=16.18$."1148 The values given in ZMR and ZRM are 7=18.50and J= 16.37. The ./ magnitude agrees fairly well. but the difference in 7 (0.16 mag) is clearly exceeding the error-bar limits.," The values given in ZMR and ZRM are $I=18.80$and $J=16.37$ The $J$ magnitude agrees fairly well, but the difference in $I$ $0.46$ mag) is clearly exceeding the error-bar limits."1149 Teidel, Teide11150to the BGL and the intergranular lane.,to the BGL and the intergranular lane.1151" Using NST TiO and observations, we discovered that small-scale intergranular jets, first described in Goodeetal.(2010b),, are associated with bright granular lanes (BGLs) developing inside photospheric granules."," Using NST TiO and observations, we discovered that small-scale intergranular jets, first described in \cite{goode_apjl_2010}, are associated with bright granular lanes (BGLs) developing inside photospheric granules."1152 The BGLs are thought to be a signature of vortex tubes and they were first found in solar and simulation data by Steineretal.(2010)., The BGLs are thought to be a signature of vortex tubes and they were first found in solar and simulation data by \citet{Steiner_2010}.1153". We summarize our new findings as follows: i) our conservative estimate is that more than half of a total of 100 well identified tiny intergranular jets are co-spatial and co-temporal with the occurrence of BGLs, although not each BGL event is accompanied with small-scale chromospheric activity; ii) along with the BGL, a vortex tube also develops a well-defined bright grain located between the BGL and the dark intergranular lane; and iii) vortex-tube signatures may reach the lower chromosphere and can be detected in off-band images."," We summarize our new findings as follows: i) our conservative estimate is that more than half of a total of 100 well identified tiny intergranular jets are co-spatial and co-temporal with the occurrence of BGLs, although not each BGL event is accompanied with small-scale chromospheric activity; ii) along with the BGL, a vortex tube also develops a well-defined bright grain located between the BGL and the dark intergranular lane; and iii) vortex-tube signatures may reach the lower chromosphere and can be detected in off-band images."1154" The bright grain, described here, appears to correspond to the plateau in the model intensity profile presented in Figure 5 in Steineretal.(2010)."," The bright grain, described here, appears to correspond to the plateau in the model intensity profile presented in Figure 5 in \cite{Steiner_2010}."1155". According to the simulation data, the darkish space between the bright grain and the BGL coincides with the axis of the vortex tube."," According to the simulation data, the darkish space between the bright grain and the BGL coincides with the axis of the vortex tube."1156" The interpretation is that due to low pressure and temperature, the opacity above the vortex tube is reduced thus allowing us to peer deeper into its relatively cooler interior."," The interpretation is that due to low pressure and temperature, the opacity above the vortex tube is reduced thus allowing us to peer deeper into its relatively cooler interior."1157 What is the bright grain then?, What is the bright grain then?1158 Is it part of the vortex tube?, Is it part of the vortex tube?1159 Does this interpretation hold when we consider the fact that the vortex tube can each the chromosphere?, Does this interpretation hold when we consider the fact that the vortex tube can reach the chromosphere?1160" As it follows from simulations, the associated magnetic field, is generally wrapped up in such a way, that the field is mainly aligned with the flow, i.e., it is rather perpendicular to the vortex tube axis."," As it follows from simulations, the associated magnetic field, is generally wrapped up in such a way, that the field is mainly aligned with the flow, i.e., it is rather perpendicular to the vortex tube axis."1161 The high-speed flow above the vortex tube reaches p to the top of the photosphere with velocities up to 8 km ! and sweeps the magnetic field in the horizontal direction to the intergranular lane., The high-speed flow above the vortex tube reaches up to the top of the photosphere with velocities up to 8 km $^{-1}$ and sweeps the magnetic field in the horizontal direction to the intergranular lane.1162" It may be that this field collides with the nearby intergranular field of possibly opposing polarity, which has the potential to cause some chromospheric activity."," It may be that this field collides with the nearby intergranular field of possibly opposing polarity, which has the potential to cause some chromospheric activity."1163" We do not know, however, if BGL events seen in images of granulation possess magnetic fields strong enough to cause detectable chromospheric activity."," We do not know, however, if BGL events seen in images of granulation possess magnetic fields strong enough to cause detectable chromospheric activity."1164 Polarization measurements with the baloon-borne solar telescope failed to detect magnetic field signal associated with a vortex tube (Steineretal.2010)., Polarization measurements with the baloon-borne solar telescope failed to detect magnetic field signal associated with a vortex tube \citep{Steiner_2010}.1165". A brief review of published Hinode/SP data (e.g.,Centenoetal.2007;2010) indicates that Hinode/SP intensity maps may have insufficient spatial resolution to reliably discern a BGL event, so that no reliable conclusions on the association between a BGL and the magnetic field can be made."," A brief review of published Hinode/SP data \citep[e.g.,][]{centeno2007,lites_2008,2010ApJ...713.1310I, gomory_2010}1166 indicates that Hinode/SP intensity maps may have insufficient spatial resolution to reliably discern a BGL event, so that no reliable conclusions on the association between a BGL and the magnetic field can be made."1167" Nevertheless, Centenoetal.(2007,leftpanelintheirFig-ure1) studied a flux-emergence event associated with a particular pattern in the intensity maps that could be interpreted as an evolving BGL."," Nevertheless, \citet[][left panel in their Figure 1]{centeno2007} studied a flux-emergence event associated with a particular pattern in the intensity maps that could be interpreted as an evolving BGL."1168" OrozcoSuárezetal.(2008) presented data for two flux-emergence events, where increased circular polarization polarization was spatial with enhanced brightness within a granule."," \cite{orozco} presented data for two flux-emergence events, where increased circular polarization polarization was co-spatial with enhanced brightness within a granule."1169 Zhangetal.(2009) reported that granules tend to fragment when magnetic fields emerge within them., \cite{Zhang_granule_fragmentation} reported that granules tend to fragment when magnetic fields emerge within them.1170 Simulations by Tortosa-Andreu&Moreno-Insertis(2009) seem to confirm the latter by showing that surface temperature structures change as field emerges., Simulations by \cite{tortosa_andreu} seem to confirm the latter by showing that surface temperature structures change as field emerges.1171" On the other hand, Gómóryetal.2010) argue that an emerging loop leaves no detectable brightness pattern on the host granule."," On the other hand, \cite{gomory_2010} argue that an emerging loop leaves no detectable brightness pattern on the host granule."1172" Stenflo(2011) underscored the possible existence of two distinct populations of the solar magnetic fields: i) strong, or collapsed, fields predominantly located in the intergranular lanes and manifested via photospheric bright points and ii) weak, or uncollapsed, flux occupying both intergranular lanes and bright granules with a wea preference for the bright granular cells."," \cite{stenflo_2011} underscored the possible existence of two distinct populations of the solar magnetic fields: i) strong, or collapsed, fields predominantly located in the intergranular lanes and manifested via photospheric bright points and ii) weak, or uncollapsed, flux occupying both intergranular lanes and bright granules with a weak preference for the bright granular cells."1173 The uncollapsed population is thought to represent weaker turbulent fields with spatial scales too small to be fully resolved with today's state-of-the-art instrumentation., The uncollapsed population is thought to represent weaker turbulent fields with spatial scales too small to be fully resolved with today's state-of-the-art instrumentation.1174" In this case, we suggest that the intergranular Jets, associated with the development of BGLs, may be a manifestation of these weaker turbulent fields, the bulk of which apparently remains hidden at spatial scales below 200 km (Stenflo2011)."," In this case, we suggest that the intergranular jets, associated with the development of BGLs, may be a manifestation of these weaker turbulent fields, the bulk of which apparently remains hidden at spatial scales below 200 km \citep{stenflo_2011}."1175. The intergranular jets are much smaller and weaker than all previously known jet-like events., The intergranular jets are much smaller and weaker than all previously known jet-like events.1176" At the same time, they appear much more numerous than the larger events, leading us to the speculation that the total energy released by these tiny events may not be negligible in the total energy balance."," At the same time, they appear much more numerous than the larger events, leading us to the speculation that the total energy released by these tiny events may not be negligible in the total energy balance."1177 Authors thank BBSO observers and the instrument team for their contribution to this study., Authors thank BBSO observers and the instrument team for their contribution to this study.1178 VY work was partly supported under NASA GI NNX08AJ20G and LWS TR&TT NNGO-5GN34G grants., VY work was partly supported under NASA GI NNX08AJ20G and LWS T NNG0-5GN34G grants.1179 VA acknowledges partial support from NSF grant ATM-0716512., VA acknowledges partial support from NSF grant ATM-0716512.1180" PG, VA and VY are partially supported by NSF (AGS-0745744), NASA (NNY 08BA22C)."," PG, VA and VY are partially supported by NSF (AGS-0745744), NASA (NNY 08BA22G)."1181 PG is partially supported by AFOSR (FA9550-09-1-0655)., PG is partially supported by AFOSR (FA9550-09-1-0655).1182 OS acknowledges insightful discussions on small-scale filament formation within the ISSI International Team lead by Ι.Ν. Kitiashvili at ISSI (International Space Science Institute) in Bern., OS acknowledges insightful discussions on small-scale filament formation within the ISSI International Team lead by I.N. Kitiashvili at ISSI (International Space Science Institute) in Bern.1183to those found in the Taurus Molecular Cloud.,to those found in the Taurus Molecular Cloud.1184 The observation of abundant ο in L1251À coustitutes the first detection of anions iu a protostcllar cuvelope outside of the Taurus-Auriga complex. aud indicates that anious are Likely to be widespread throughout Calactic star-forming regions where carbon chains are present.," The observation of abundant $_6$ $^-$ in L1251A constitutes the first detection of anions in a protostellar envelope outside of the Taurus-Auriga complex, and indicates that anions are likely to be widespread throughout Galactic star-forming regions where carbon chains are present."1185 The techuique of using UC3N aud CT as proxies for the detection of less abundaut. larger carbon chains aud anions shows strong promise as a leans for obtainiug further detections of these inolecules iu low-inass forming regions in the future.," The technique of using $_3$ N and $_4$ H as proxies for the detection of less abundant, larger carbon chains and anions shows strong promise as a means for obtaining further detections of these molecules in low-mass star-forming regions in the future."1186 This researcli was supported by the NASA Exobicloey Program and the Coddard Center for Astrobioloey., This research was supported by the NASA Exobiology Program and the Goddard Center for Astrobiology.1187 Astrophysics at QUB is supported by a eraut from STFC., Astrophysics at QUB is supported by a grant from STFC.1188Karl-Sehhwarzschild-Str.2.85748labelsect:intro Variability is a dominant observational signature in pre-main sequence (pre-MS) stellar evolution.,"hwarzschild-Str.\tikzmark{mainBodyEnd0} \tikzmark{mainBodyStart1}2,\tikzmark{mainBodyEnd1} \tikzmark{mainBodyStart2}85748\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}Garching,\tikzmark{mainBodyEnd3} \tikzmark{mainBodyStart4}Germany\tikzmark{mainBodyEnd4} \tikzmark{mainBodyStart5}}\tikzmark{mainBodyEnd5} 1189 1190%\titlerunning{X-ray emission from Z\,CMa during outburst and from its jet}1191 1192\date{Received $<$29-01-2009$>$ / Accepted $<$16-03-2009$>$} 1193 1194 1195\abstract{1196Accretion shocks have been recognized as important X-ray emission mechanism 1197for pre-main sequence stars. Yet the X-ray properties of FUor outbursts, events that are caused by1198violent accretion, have been given little attention. We have observed the FUor1199object Z\,CMa during optical outburst and quiescence with {\em Chandra}. No significant1200changes in X-ray brightness and spectral shape are found, suggesting that the X-ray 

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