ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 At small optical depth. Gxf. whereas at high optical depth. the function is given by GzI.," At small optical depth, $G \propto t$, whereas at high optical depth, the function is given by $G\approx 1$."3 This figure shows how cTuas of differeut optical depths will contribute iu a luge euseuble., This figure shows how clumps of different optical depths will contribute in a large ensemble.4" However. 1cting that fXf?goMfF2.11 Fiewre 9 also sugeests how a single chip contributes o the euseiuble teMal at different frequencics,"," However, noting that $t \propto5f^{-2} \,g_\nu \propto f^{-2.11}$, Figure \ref{fig2} also suggests how a single clump contributes to the ensemble total at different frequencies."6 From this perspective low f£ corresponds to high Εν and vice scorsa.," From this perspective low $t$ corresponds to high $f$, and vice versa."7The implicaion is that for uv eiven cTu» there will always be frequencies at which t1C ‘hump is thii and frequencies where it is thick.,"The implication is that for any given clump, there will always be frequencies at which the clump is thin and frequencies where it is thick."8 Thus οςuation (7)) is found to have two iutant Bits., Thus equation \ref{eq:radioflux}) ) is found to have two important limits.9 For clumps that are very thick (1co. fI). je resultan SED will be proportional to f?Gxf?.," For clumps that are very thick (i.e., $t \gg 1$ ), the resultant SED will be proportional to $f^2\,G \propto f^2$."10 I£ all the chuups are thin. then the SED wi] approach. m2xf0. which is the scaling for the Cant factor in the radio band.," If all the clumps are thin, then the SED will approach $f^2\,G \propto f^{-0.11}$ which is the scaling for the Gaunt factor in the radio band."11 The variation betweei thick aud in thus eies a power-law range between 0.11 andl 2 so that in principle. one can inaeiue a frequency “window over which the SED may approximae a power-law with a slope that is itermeciae between trese two limiting values.," The variation between thick and thin thus gives a power-law range between $-0.11$ and $+2$, so that in principle, one can imagine a frequency “window” over which the SED may approximate a power-law with a slope that is intermediate between these two limiting values."12 Note that we are ignoring contributions from the central star because the rac10 continuuni enussion of the star will be dwarfed w the coutribution frou the ΠΤΙ region: howey. the star and its immediate circumstellar component may start to compete with the exteuded emission at he shorter nuu and far-IR wavelengths.," Note that we are ignoring contributions from the central star because the radio continuum emission of the star will be dwarfed by the contribution from the HII region; however, the star and its immediate circumstellar component may start to compete with the extended emission at the shorter mm and far-IR wavelengths."13 We now introduce a power-law nunuber distribution iu terms of the optical depth along the chup Haneter: such that thie total πανο of clumps Vey is given by and Ny- is. a normalization. constant., We now introduce a power-law number distribution in terms of the optical depth along the clump diameter: such that the total number of clumps $N_{\rm cl}$ is given by and $N_0$ is a normalization constant.14 Now- mean physical. parameters such as (Dj/ aud (ή)iP are defined with reference to this optical depth distribution., Now mean physical parameters such as $\langle T \rangle$ and $\langle R^2 \rangle$ are defined with reference to this optical depth distribution.15 Care must be taken im handling this distribution. since f=f( f). ," Care must be taken in handling this distribution, since $t = t(f)$ "16"by 2.. and a saturated value of L7""/Li4=1040 js used. which is the average of the highest mass group in Pizzolato's work.","by \cite{pizzolato}, and a saturated value of $L_{x}^{sat}/L_{bol} = 10^{-3.9}$ is used, which is the average of the highest mass group in Pizzolato's work."17 A difference between our analysis and. the population study of 7. is that Penz ct al., A difference between our analysis and the population study of \cite{penzXevap} is that Penz et al.18 included: a range of. X- behaviour for stars of cach type. whereas we assume average properties. since we do not know the individual histories of the objects in our study.," included a range of X-ray behaviour for stars of each type, whereas we assume average properties, since we do not know the individual histories of the objects in our study."19 Our estimates of N-rav irradiation would be biassed if the X-ray emission. of stars were correlated with the presence of planets., Our estimates of X-ray irradiation would be biassed if the X-ray emission of stars were correlated with the presence of planets.20 Indeed. ? [ind evidence for excess X-ray emission [rom stars with known planets.," Indeed, \cite{excessXRAY} find evidence for excess X-ray emission from stars with known planets."21 We choose not to moclily our estimates by this factor because we are interested. in the duration and strength of the saturated phase of X-ray emission. whereas the measurements of Ixasvap ct aapply to the weaker X-rav emission of much older stars.," We choose not to modify our estimates by this factor because we are interested in the duration and strength of the saturated phase of X-ray emission, whereas the measurements of Kasyap et apply to the weaker X-ray emission of much older stars."22 Nevertheless. if the results of 7. were found to apply to stars of all ages. then we may have underestimated the X-ray irradiation of our sample by around half an order of magnitiude.," Nevertheless, if the results of \cite{excessXRAY} were found to apply to stars of all ages, then we may have underestimated the X-ray irradiation of our sample by around half an order of magnitiude."23 ligure 2 compares the potential energy. of planets in our sample with the energy incident during the saturated X-ray emission phase of their parent star., Figure \ref{sat} compares the potential energy of planets in our sample with the energy incident during the saturated X-ray emission phase of their parent star.24 As can be clearly seen. all of the systems considered in this work are able to survive the energy absorbed during the saturation period of their parent star.," As can be clearly seen, all of the systems considered in this work are able to survive the energy absorbed during the saturation period of their parent star."25 Of course. each of these svstenis has already weathered one saturation period. so Fig.," Of course, each of these systems has already weathered one saturation period, so Fig."26 2. really tests whether these systems could survive a second saturation period., \ref{sat} really tests whether these systems could survive a second saturation period.27 Since they all could survive another saturation period. it seems likely that they were only minimally alfected by the first. period.," Since they all could survive another saturation period, it seems likely that they were only minimally affected by the first period."28In order to convert the luminosity density profiles to stellar mass density profiles we make the following assumptions about the mass-to-light (M/L) ratios.,In order to convert the luminosity density profiles to stellar mass density profiles we make the following assumptions about the mass-to-light $M/L$ ) ratios.29 For the high redshift sample we use stellar masses fromKriek (2008) adjusted to aKroupa (2001) IMF., For the high redshift sample we use stellar masses from (2008) adjusted to a (2001) IMF.30" For the nearby sample, we use the well-established relation between M/L ratio and mass to convert luminosities to masses (e.g., 1991)."," For the nearby sample, we use the well-established relation between $M/L$ ratio and mass to convert luminosities to masses (e.g., 1991)."31" The normalization and slope of the relation in the B band were determined by combining the information in Table 1 ofMarel (1991) and Table 2 in (2007): The conversion from luminosity to mass is the largest in the methodology, in particular the lack of dynamical uncertaintymeasurements that could calibrate the M/L ratios of the high redshift galaxies."," The normalization and slope of the relation in the $B$ band were determined by combining the information in Table 1 of (1991) and Table 2 in (2007): The conversion from luminosity to mass is the largest uncertainty in the methodology, in particular the lack of dynamical measurements that could calibrate the $M/L$ ratios of the high redshift galaxies."32" We will return to this issue in ,66.", We will return to this issue in 6.33 The stellar density profiles of the compact high redshift galaxies are compared to those of nearby elliptical galaxies in reffig:densprof.., The stellar density profiles of the compact high redshift galaxies are compared to those of nearby elliptical galaxies in \\ref{fig:densprof}.34 The solid line is the average density profile of the 9 galaxies from vD08., The solid line is the average density profile of the 9 galaxies from vD08.35 We use a 1000 iteration bootstrap estimation to approximate errors of the profile due to the small sample size of the high averageredshift densitygalaxies., We use a 1000 iteration bootstrap estimation to approximate errors of the average density profile due to the small sample size of the high redshift galaxies.36 The | o contour is shown in dark gray and the 2 o is shown in light gray., The 1 $\sigma$ contour is shown in dark gray and the 2 $\sigma$ is shown in light gray.37 Broken lines are average profiles of nearby from the Τ09 in three different mass ellipticalbins.," Broken lines are average profiles of nearby elliptical galaxies from the T09 sample, in three different mass bins."38" The galaxieslowest mass bin is sample,M>10!Ma: this is the mass that the high redshift galaxies already have at the epoch of observation, and therefore the minimum mass of their descendants."," The lowest mass bin is $M \ge 10^{11}M_{\odot}$: this is the mass that the high redshift galaxies already have at the epoch of observation, and therefore the minimum mass of their descendants."39 Figure 1 shows that the discrepancy between the profiles of compact high redshift galaxies and nearby elliptical galaxies is mostly in the outer regions., Figure \ref{fig:densprof} shows that the discrepancy between the profiles of compact high redshift galaxies and nearby elliptical galaxies is mostly in the outer regions.40" Within r+1kpc the average stellar density of the high redshift galaxies is greater than the density of nearby ellipticals by a factor of a few only, particularly for the more massive galaxies in the TOO sample."," Within $r \approx 1\unit{kpc}$ the average stellar density of the high redshift galaxies is greater than the density of nearby ellipticals by a factor of a few only, particularly for the more massive galaxies in the T09 sample."41" This is much smaller than the factor of =100 difference discrepancywhen the density is measured within the effective radius (e.g., vD08)."," This discrepancy is much smaller than the factor of $\gtrsim 100$ difference when the density is measured within the effective radius (e.g., vD08)."42" Furthermore, our error estimates only address the sample bias; this discrepancy is especially small considering the other sources of uncertainty in our measurements, which we will discuss further in ,66."," Furthermore, our error estimates only address the sample bias; this discrepancy is especially small considering the other sources of uncertainty in our measurements, which we will discuss further in 6."43" Outside of this inner region, the difference grows significantly; the stellar density of nearby elliptical galaxies is a few hundred times higher than that of the compact high redshift galaxies at r lOkpc."," Outside of this inner region, the difference grows significantly; the stellar density of nearby elliptical galaxies is a few hundred times higher than that of the compact high redshift galaxies at $r>10 \unit{kpc}$."44" We infer that in order to evolve into nearby galaxies, compact galaxies need not significantly in the central regions, but must grow changesignificantly in their outer regions."," We infer that in order to evolve into nearby galaxies, compact galaxies need not change significantly in the central regions, but must grow significantly in their outer regions."45 The relative properties of the high redshift galaxies and nearby galaxies are demonstrated in Figure 2.., The relative properties of the high redshift galaxies and nearby galaxies are demonstrated in Figure \ref{fig:allprop}.46 The compact galaxies from vD08 are indicated by solid circles., The compact galaxies from vD08 are indicated by solid circles.47 The nearby samples are represented by open symbols: squares for TOO and triangles for FPJ., The nearby samples are represented by open symbols: squares for T09 and triangles for FPJ.48" Only nearby galaxies with sufficient mass, M>10!Mo, can be the descendants of the high redshift galaxies."," Only nearby galaxies with sufficient mass, $M \gtrsim 10^{11} M_{\odot}$, can be the descendants of the high redshift galaxies."49 Galaxies with lower masses are denoted with light grey symbols., Galaxies with lower masses are denoted with light grey symbols.50 The relative compactness of high redshift and low redshift galaxies is shown in Figure 2((a)., The relative compactness of high redshift and low redshift galaxies is shown in Figure \ref{fig:allprop}( (a).51 There is a clear trend showing the increasing effective radius with galaxy mass in the nearby galaxies., There is a clear trend showing the increasing effective radius with galaxy mass in the nearby galaxies.52" The high redshift galaxies, though in the middle of the nearby mass range, are smaller by a factor of ~5 in effective radius."," The high redshift galaxies, though in the middle of the nearby mass range, are smaller by a factor of $\sim 5$ in effective radius."53" This result confirms previous studies, which generally used the Sloan Digital Sky Survey (SDSS) as a low redshift comparison point (e.g.,Toft 2007; 2008; 2008; 2008). ("," This result confirms previous studies, which generally used the Sloan Digital Sky Survey (SDSS) as a low redshift comparison point (e.g., 2007; 2008; 2008; 2008). ("54"2008) combine data from the literature (in addition to adding new data at z= 1) and derive an evolution of r,ox(1--z)-1299? at fixed mass (for samples with photometrically determined masses), corresponding to a factor of 4.2t0.6 at z22.3.","2008) combine data from the literature (in addition to adding new data at $z\approx 1$ ) and derive an evolution of $r_e55\propto (1+z)^{-1.20\pm 0.12}$ at fixed mass (for samples with photometrically determined masses), corresponding to a factor of $4.2\pm 0.6$ at $z=2.3$."56 The difference in size at fixed mass implies a significant difference in density contained within the effective radius of the high redshift and nearby galaxies., The difference in size at fixed mass implies a significant difference in density contained within the effective radius of the high redshift and nearby galaxies.57" We calculated the average densities within the effective radius by integrating the stellar density profiles derived in the previous Section: with r=r,.", We calculated the average densities within the effective radius by integrating the stellar density profiles derived in the previous Section: with $r=r_e$.58 This difference is obvious in Figure 2((b): the vertical axis of this panel demonstrates the factor of >100 differences in the average density within the effective radius., This difference is obvious in Figure \ref{fig:allprop}( (b): the vertical axis of this panel demonstrates the factor of $>100$ differences in the average density within the effective radius.59" The horizontal axis of reffig:allprop((b) showsthe average density integrated to kkpc rather than r= r,.", The horizontal axis of \\ref{fig:allprop}( (b) showsthe average density integrated to $r=1$ kpc rather than $r=r_e$ .60" For convenience, we will refer to"," For convenience, we will refer to"61(Lyubimkov et al.,(Lyubimkov et al.62 20043., 2004).63 When adopting. for instance. ¢=400à. one may find from Heger Langer's (2000) computations that fora 15 Mo star the surface C. N and O abundances change by -|.2. 40.9 and -0.4 dex. respectively.," When adopting, for instance, $v=400$, one may find from Heger Langer's (2000) computations that for a 15 Mo star the surface C, N and O abundances change by -1.2, +0.9 and -0.4 dex, respectively."64 Therefore. use of C. N and O as indicators of metallicity of early B stars is not strictly correct.," Therefore, use of C, N and O as indicators of metallicity of early B stars is not strictly correct."65" In the case of cool stars. iron is usually used as an indicator of metallicity,"," In the case of cool stars, iron is usually used as an indicator of metallicity."66 On the one hand. this element does not participate in thermonuclear reactions during early evolutionary phases. and. therefore. its abundance corresponds to the initial metallicity of a star.," On the one hand, this element does not participate in thermonuclear reactions during early evolutionary phases, and, therefore, its abundance corresponds to the initial metallicity of a star."67 On the other hand. it gives a lot of spectral lines. which allow its abundance to be derived with high accuracy.," On the other hand, it gives a lot of spectral lines, which allow its abundance to be derived with high accuracy."68 The metallicity ΠΟΠΗ] is determined then as the difference between the Fe abundances of the star and the Sun. i.e. where the abundance log= corresponds to the standard logarithmic scale. with the hydrogen abundance logz(ll)=19.00.," The metallicity $Fe/H$ ] is determined then as the difference between the Fe abundances of the star and the Sun, i.e. where the abundance $\log \varepsilon$ corresponds to the standard logarithmic scale with the hydrogen abundance $\log69\varepsilon({\rm H})=12.00$."70 Unfortunately. in spectra of early B stars the iron lines are weak.," Unfortunately, in spectra of early B stars the iron lines are weak."71 One of the strongest lines in spectra of early and medium B stars is the AHS1.2 lline., One of the strongest lines in spectra of early and medium B stars is the 4481.2 line.72 Magnesium is quite suitable as an indicator of metallicity. because it does not change its abundance during the MS evolutionary phase.," Magnesium is quite suitable as an indicator of metallicity, because it does not change its abundance during the MS evolutionary phase."73 Even its participation in the MgAl-cycle leads to a negligible decrease of the abundance. by about 0.02—0.03 dex (Daflon et al.," Even its participation in the MgAl-cycle leads to a negligible decrease of the abundance, by about 0.02–0.03 dex (Daflon et al."74 2003)., 2003).75 Moreover. for magnesium. unlike C. N and O. the precise solar abundance is known.," Moreover, for magnesium, unlike C, N and O, the precise solar abundance is known."76 In fact. according to Holweger (20013. its abundance derived from photospheric lines is logz(Mg)=7.54+0.06.," In fact, according to Holweger (2001), its abundance derived from photospheric lines is $\log \varepsilon({\rm77Mg})=7.54\pm0.06$."78 This spectroscopic value is very close to the meteoritic abundance 3:7.560.02., This spectroscopic value is very close to the meteoritic abundance $7.56\pm0.02$.79 Averaging these two abundances. Lodders (2003) recommends the value logz.;(Mg)=1.55dE0.02 as the solar magnesium abundance (meteoritic’ refers to the CI-type chondrites).," Averaging these two abundances, Lodders (2003) recommends the value $\log80\varepsilon_{\sun}({\rm Mg})= 7.55\pm0.02$ as the solar magnesium abundance ('meteoritic' refers to the CI-type chondrites)."81 We shall use the latter value., We shall use the latter value.82" The metallicity of stars will be determined from formula It should be noted that Lodders presents also the proto-Sun magnesium abundance logs,.(Alg)=7.62+0.02."," The metallicity of stars will be determined from formula It should be noted that Lodders presents also the proto-Sun magnesium abundance $\log \varepsilon_{ps}({\rm83Mg})=7.62\pm0.02$."84 This value differs slightly from the above-mentioned photospheric value because heavy-element fractionation in the Sun has altered original abundances., This value differs slightly from the above-mentioned photospheric value because heavy-element fractionation in the Sun has altered original abundances.85 At present we continue our study of large sample of early and medium B stars., At present we continue our study of large sample of early and medium B stars.86" Our general goal was to provide new comprehensive observations of the stars and to determine and analyse the abundances of CNO-cycle elements. namely helium. carbon, nitrogen and oxygen."," Our general goal was to provide new comprehensive observations of the stars and to determine and analyse the abundances of CNO-cycle elements, namely helium, carbon, nitrogen and oxygen."87 A series of papers is being published., A series of papers is being published.88 First. we obtained high-resolution spectra of more than 100 stars at two observatories. namely the McDonald Observatory of the University of Texas and the Crimean Astrophysical Observatory (Ukraine): equivalent widths of 2 hydrogen lines and 11 helium ines were measured (Lyubimkov et al.," First, we obtained high-resolution spectra of more than 100 stars at two observatories, namely the McDonald Observatory of the University of Texas and the Crimean Astrophysical Observatory (Ukraine); equivalent widths of 2 hydrogen lines and 11 helium lines were measured (Lyubimkov et al."89 2000. hereinafter Paper D.," 2000, hereinafter Paper I)."90" Second. we determined a number of fundamental parameters of he stars including the effective temperature7;4r.. surface gravity g.. interstellar extinction ;1,. distance d. mass AJ. radius /?. uminosity £L. age /. and relative age///sjs.. where fare is the MS ifetime (Lyubimkov et al."," Second, we determined a number of fundamental parameters of the stars including the effective temperature, surface gravity , interstellar extinction $A_v$ , distance $d$, mass $M$ radius $R$, luminosity $L$, age $t$, and relative age, where $t_{\rm91MS}$ is the MS lifetime (Lyubimkov et al."92 2002. Paper ID.," 2002, Paper II)."93 Third. from non-LTE analysis of lines we derived the helium abundance //0/1/. microturbulent yarameter aand projected rotational velocity (Lyubimkov et al.," Third, from non-LTE analysis of lines we derived the helium abundance $He/H$, microturbulent parameter and projected rotational velocity (Lyubimkov et al."94 2004. Paper IIT).," 2004, Paper III)."95 Presently.7 we are analysingCi. aand lines: in particular. the microturbulent parameter iis determined from these lines. too.," Presently, we are analysing, and lines; in particular, the microturbulent parameter is determined from these lines, too."96 We present in this paper a new determination of the Mg abundances in B stars that is based on our spectra and our atmospheric parameters., We present in this paper a new determination of the Mg abundances in B stars that is based on our spectra and our atmospheric parameters.97 We use the A4481.2 lline and for a sample of hot stars a weaker line at 7877Α., We use the 4481.2 line and for a sample of hot stars a weaker line at 7877.98. An analysis of the LHH8I.2À lline can be complicated by presence of the neighbouring 44479.9 line., An analysis of the 4481.2 line can be complicated by presence of the neighbouring 4479.9 line.99 Tf a star has a rather high rotational velocity¢sin/.. the two lines form a common blend.," If a star has a rather high rotational velocity, the two lines form a common blend."100 In such cases the computation of synthetic spectra is necessary., In such cases the computation of synthetic spectra is necessary.101 The 4479.9 line intensity depends strongly onYiu., The 4479.9 line intensity depends strongly on.102" In Fig.l we show as examples the observed spectra in the A4481.2 vicinity for three B stars with the different effective temperatures aand low rotational velocities (νο, with sharp spectral lines).", In Fig.1 we show as examples the observed spectra in the 4481.2 vicinity for three B stars with the different effective temperatures and low rotational velocities (i.e. with sharp spectral lines).103 One sees that the 44479.9 lline is relatively strong at == 22500 K. but weak both at the high temperature. == 30700 K. and atthe lowone.== 15500 K.," One sees that the 4479.9 line is relatively strong at = 22500 K, but weak both at the high temperature, = 30700 K, and atthe lowone,= 15800 K."104A similar result applies to simulations where :7. is reduced and then held at some value δρ (such as 20).,A similar result applies to simulations where $\beta_c$ is reduced and then held at some value $\beta_{hold}$ (such as \citealt{Clarke2007}) ).105 This action will fix Ly at some value given by the mismatch between the first and second terms in equation (23))., This action will fix $\Gamma_J$ at some value given by the mismatch between the first and second terms in equation \ref{eq:betaj}) ).106" Low values of ια will increase the mismatch. decreasing Ly and reducing the timescale or fragmentation to occur,"," Low values of $\beta_{hold}$ will increase the mismatch, decreasing $\Gamma_J$ and reducing the timescale for fragmentation to occur."107 Given this and equation (339). we can now make some general statements about the work of ?..," Given this and equation \ref{eq:dbetajdy}) ), we can now make some general statements about the work of \citet{Clarke2007}. ."108" They see some dises fragmenting at very low 3. because once ‘avourable fragmentation conditions arise. there is a finite timescale ‘or fragmentation to occur. defined by Ε.Π 7 is sufficiently small (to the point of being unphysical). then ;2.. can become extremely small before LF), reaches values necessary for fragmentation."," They see some discs fragmenting at very low $\beta_c$ because once favourable fragmentation conditions arise, there is a finite timescale for fragmentation to occur, defined by $\Gamma_J$.If $T$ is sufficiently small (to the point of being unphysical), then $\beta_c$ can become extremely small before $\Gamma_J$ reaches values necessary for fragmentation."109" Also. we show that as ? note themselves. if they were able to run their high 7 simulations with .3,. below the critical value for a sufficiently long time. they would also probably fragment."," Also, we show that as \citet{Clarke2007} note themselves, if they were able to run their high $T$ simulations with $\beta_c$ below the critical value for a sufficiently long time, they would also probably fragment."110 The length of simulation time required to produce a fragment will be calculable from the evolution of L;;., The length of simulation time required to produce a fragment will be calculable from the evolution of $\Gamma_J$.111 While high 7 simulations may be more stable to fragmentation in the short-term. their long-term stability is not likely to be much better than simulations where the cooling time is rapidly changing.," While high $T$ simulations may be more stable to fragmentation in the short-term, their long-term stability is not likely to be much better than simulations where the cooling time is rapidly changing."112 As a final aside on the subject. we must be careful about the sign of s.," As a final aside on the subject, we must be careful about the sign of $\dot{\beta_c}$."113" Fragmentation is more likely if DL; is negative. and hence 3, should be negative also."," Fragmentation is more likely if $\Gamma_J$ is negative, and hence $\dot{\beta}_c$ should be negative also."114" If the change in cooling time is rapid andpositive.. then 2, will become too large. and the Jeans mass will either increase or become steady at a higher value (which can be seen from the Pa;7. dependence of equation (17)))."," If the change in cooling time is rapid and, then $\beta_c$ will become too large, and the Jeans mass will either increase or become steady at a higher value (which can be seen from the $\beta_c$ dependence of equation \ref{eq:mjeans_sigma}) ))."115". Large fluctuations in «2, are therefore not a guarantee of fragmentation -favourable.", Large fluctuations in $\beta_c$ are therefore not a guarantee of fragmentation -.116. This is akin to the fragmentation of shocked flows producing star formation in molecular clouds (e.g. 229). and may explain the formation of objects during protostellar encounters with sufficiently extended dises (222)..," This is akin to the fragmentation of shocked flows producing star formation in molecular clouds (e.g. \citealt{Heitsch2008,Bonnell2008}) ), and may explain the formation of objects during protostellar encounters with sufficiently extended discs \citep{Thies2005,Shen2010,Thies2010}."117 Equally. this accounts for Tagmentation suppression during encounters with compact dises (22)... where the subsequent increase in optical depth ensures 1. is arge and positive.," Equally, this accounts for fragmentation suppression during encounters with compact discs \citep{Lodato_encounters,encounters}, where the subsequent increase in optical depth ensures $\dot{\beta_c}$ is large and positive."118" Let us assume for the moment that heating and cooling are in exact balance (with a large. constant value of .3,.). such that the heating and cooling terms in equation 25 cancel and «3.=0. and hence we can approximate i.e. glibly speaking. dises which have a higher accretion rate are more favourable to fragmentation."," Let us assume for the moment that heating and cooling are in exact balance (with a large, constant value of $\beta_c$ ), such that the heating and cooling terms in equation \ref{eq:betaj} cancel and $\dot{\beta}_c=0$, and hence we can approximate i.e. glibly speaking, discs which have a higher accretion rate are more favourable to fragmentation."119 Of course. accretion of envelope material does not merely increase the dise mass - it affects the local thermodynamics and the angular momentum distribution. both of which ean push the dise away from (2.— 1.," Of course, accretion of envelope material does not merely increase the disc mass - it affects the local thermodynamics and the angular momentum distribution, both of which can push the disc away from $Q\sim 1$ ."120 We should therefore also demand if the dise is to remain amenable to fragmentation., We should therefore also demand if the disc is to remain amenable to fragmentation.121 Substituting de;fd=&./™ Cand similarly for (2). we obtain Upon rearrangement. the following condition now appears: At the critical point(where M=(QJ this can be simplified to merely If accretion results in greater deposition of angular momentum or in local heating. then the dise moves away from fragmentation (this is intuitively obvious from the detinition of (9).," Substituting $d c_s / d\Sigma = \dot{c}_s/ \dot{\Sigma}$ (and similarly for $\Omega$ ), we obtain Upon rearrangement, the following condition now appears: At the critical point(where $\frac{dQ}{d\Sigma}=0$ ), this can be simplified to merely If accretion results in greater deposition of angular momentum or in local heating, then the disc moves away from fragmentation (this is intuitively obvious from the definition of $Q$ )."122 Also. it further underlines the stabilising influence of sustained heating. either by accretion shocks or by stellar irradiation (222). ," Also, it further underlines the stabilising influence of sustained heating, either by accretion shocks or by stellar irradiation \citep{Rafikov_07,Mejia_4,Stamatellos2008}. ."123"Marginally stable. self-gravitating discs are susceptible to fragmentation if: We should be able to identify an equilibrium Jeans Mass by demanding A;=0. or This is a first order non-linear ODE for :7,. (assuming a constant Is "," Marginally stable self-gravitating discs are susceptible to fragmentation if: We should be able to identify an equilibrium Jeans Mass by demanding $\dot{M}_J=0 $, or This is a first order non-linear ODE for $\beta_c$ (assuming a constant $\Gamma_\Sigma$ )."124In the constant «3. formalisms of the past. this becomes These formalisms are not predictors of the fragment massse.. but thev show the secular Jeans mass. and to some extent the secular fragment mass if the disc is left to its own devices and satisfies the correct criteria.," In the constant $\beta_c$ formalisms of the past, this becomes These formalisms are not predictors of the fragment mass, but they show the secular Jeans mass, and to some extent the secular fragment mass if the disc is left to its own devices and satisfies the correct criteria."125 Also. in practice the cooling rate will be a function of deposition rate (in the sense that this maintains a local optical depth). the angular velocity will be a function of mass ratio and the cooling rate will be sensitive to changes in temperature and surface density.," Also, in practice the cooling rate will be a function of deposition rate (in the sense that this maintains a local optical depth), the angular velocity will be a function of mass ratio and the cooling rate will be sensitive to changes in temperature and surface density."126 This differential equation is only soluble numerically. and probably only with numerical simulations.," This differential equation is only soluble numerically, and probably only with numerical simulations."127 In the ease of simple analytic models. it is more expedient to impose the fragmentation criteria. and calculate A7; at the fragmentation boundary.," In the case of simple analytic models, it is more expedient to impose the fragmentation criteria, and calculate $M_J$ at the fragmentation boundary."128 This method will now be applied to such dise models in the following sections., This method will now be applied to such disc models in the following sections.129 To investigate the dependence of theJeans mass on dise parameters. we construct a model of a self-gravitating dise assuming a steady state accretion rate 1. and an outer disc radius ron.," To investigate the dependence of theJeans mass on disc parameters, we construct a model of a self-gravitating disc assuming a steady state accretion rate $\dot{M}$ and an outer disc radius $r_{\rm out}$."130 We take an approach essentially identical to that of 2.. ? and ?..," We take an approach essentially identical to that of \citet{Levin2007}, \citet{Clarke_09} and \citet{Rice_and_Armitage_09}."131" For the sake of simplicity. weassume X— 0. 5, is small. and impose ο<Ly as the criterion for fragmentation."," For the sake of simplicity, weassume $\dot{\Sigma}=0$ , $\dot{\beta_c}$ is small, and impose $-5<\Gamma_J <0$ as the criterion for fragmentation."132" Cautioned by the results of ?.. we impose a saturatingvalue fora=as, 0.1."," Cautioned by the results of \citet{Cossins2010},, we impose a saturatingvalue for $\alpha = \alpha_{sat}=0.1$ ."133 The model is constructed as follows: we assume the dise is Keplerian and gravitationally unstable at all radii. Le.," The model is constructed as follows: we assume the disc is Keplerian and gravitationally unstable at all radii, i.e."134cent occurs at sub-LEdclington accretion rates. and a relevant raction is absorbed ancl re-radiated at longer wavelengths. ρα{ισατν in the infrared.,"cent occurs at sub-Eddington accretion rates, and a relevant fraction is absorbed and re-radiated at longer wavelengths, particularly in the infrared."135 This means that most of the sky as seen from an average active nucleus is obscured., This means that most of the sky as seen from an average active nucleus is obscured.136 Aloreover. given the requirement that these low uminosity AGN are at least one order of magnitude fainter han typical broad line objects (see Section 2). we predict hat starburst obscured ACN may outnumber brighter. unobseured AGN by a factor of LOO or more if they are o account for the remainder of the NRB.," Moreover, given the requirement that these low luminosity AGN are at least one order of magnitude fainter than typical broad line objects (see Section 2), we predict that starburst obscured AGN may outnumber brighter, unobscured AGN by a factor of 100 or more if they are to account for the remainder of the XRB."137 This estimate is in good agreement. with the number count. predictions or QSOs and NLXCGs based on extrapolating deep ROSAT observations (Almaini Fabian 1997), This estimate is in good agreement with the number count predictions for QSOs and NLXGs based on extrapolating deep ROSAT observations (Almaini Fabian 1997).138 If most of accretion in the Universe is highly: obseurect. hen the amount of emitted: power per unit. galaxy based on optical or UV. QSO luminosity functions (Soltan 1982. hinney 1997). and. therefore the mass in. black holes in AGN. might have been underestimatecl," If most of accretion in the Universe is highly obscured, then the amount of emitted power per unit galaxy based on optical or UV QSO luminosity functions tan 1982, Phinney 1997), and therefore the mass in black holes in AGN, might have been underestimated."139 Ελπίς is due to he fact that 90 per cent. of the accretion power would oe Obscured and. re-racdiated in the infrared., This is due to the fact that 90 per cent of the accretion power would be obscured and re-radiated in the infrared.140 Also. if most AGN are Advection Dominated Accretion Flows (ADAFs). as it has been proposed by Di Matteo Fabian (1997) as he source of the ΧΙΟ. the implied. low mass to energy conversion clliciency also means that the black hole masses would have to be larger. in agreement with the local estimates.," Also, if most AGN are Advection Dominated Accretion Flows (ADAFs), as it has been proposed by Di Matteo Fabian (1997) as the source of the XRB, the implied low mass to energy conversion efficiency also means that the black hole masses would have to be larger, in agreement with the local estimates."141 We thank Steinn Sigurdsson for very interesting comments., We thank Steinn Sigurdsson for very interesting comments.142 ACE thanks the Roval Society for support., ACF thanks the Royal Society for support.143 XB acknowledges partial financial support. provided. by the DCHES uncer project PB95-0122 and funding for his sabbatical at Cambridge under DGES grant. PRO5-490., XB acknowledges partial financial support provided by the DGES under project PB95-0122 and funding for his sabbatical at Cambridge under DGES grant PR95-490.144magnetic field.,magnetic field.145 Equation (1) is the general expression for perturbative treatment., Equation (1) is the general expression for perturbative treatment.146 It is pointed out bv Fleishman(2006a).. (he rectilinear motion of electron is valid lor laree frequencies. however. al low frequencies. the particle trajectory. traverses several correlation lengths scattering by magnetic inhomogeneities. thus the particle deflection angle accumulated along the coherence length exceeds the beaming angle (see Figure 1of Fleishman(2006a))).," It is pointed out by \citet{fleishman06a}, the rectilinear motion of electron is valid for large frequencies, however, at low frequencies, the particle trajectory traverses several correlation lengths scattering by magnetic inhomogeneities, thus the particle deflection angle accumulated along the coherence length exceeds the beaming angle (see Figure 1of \citet{fleishman06a}) )."147 The dispersion relation qo=qu(q) of the non-relativistic plasma was presented in (1959)., The dispersion relation $q_0=q_0(q)$ of the non-relativistic plasma was presented in \citet{weibel59}.148. The improved equations for the isotropic and relativistic plasma were given in detail by Mikliailovski(1980). while Yoon&Davidson(1987) built the analvtical model [or the relativistic plasma wilh a waterbag distribution., The improved equations for the isotropic and relativistic plasma were given in detail by \citet{mikhailovski80} while \citet{yoon87} built the analytical model for the relativistic plasma with a waterbag distribution.149 More comprehensive works have been performed recently by Silvaοἱal.(2002).. Wiersma&Achterberg(2004) and (2006).," More comprehensive works have been performed recently by \citet{silva02}, \citet{wiersma04} and \citet{fiore06}."150. In this paper. we choose the dispersion relation of relativistic collisionless shocks considered by Milosavljevic.Nakar&Spitkovsky(2006).," In this paper, we choose the dispersion relation of relativistic collisionless shocks considered by \citet{milosavlj06}."151. Weibel instability is an efficient wav to generate the random magnetic field in relativistic shocks (Silvaetal.2003:Schlickeiser&ShuklaWiersmaAchterberg2004:butseealsoLyubarsky&Eichler 2006).," Weibel instability is an efficient way to generate the random magnetic field in relativistic shocks \citep{silva03,schlickeiser03,wiersma04,lyubarsky06}."152. However. there could be other possibilities to form magnetic structure.," However, there could be other possibilities to form magnetic structure."153 In this Letter. we argue that the local and random magnetic field generated. by (turbulence is also relevant for the perturbative DSR. theory.," In this Letter, we argue that the local and random magnetic field generated by turbulence is also relevant for the perturbative DSR theory."154 The spectrum energv in a fully developed. turbulent. {hud can be described by the ]xolnogorov form with (he classical index —5/3., The spectrum energy in a fully developed turbulent fluid can be described by the Kolmogorov form with the classical index $-5/3$.155 For the magnetic turbulence. the cascade delay time max enter (he estimation of energy (rausler rate. (he energv spectrum of Ixraichinan has an index of —3/2.," For the magnetic turbulence, the cascade delay time may enter the estimation of energy transfer rate, the energy spectrum of Kraichnan has an index of $-3/2$."156 Although the situation we focus on has no external magnetic field. ab small scales. the turbulence is still shown as the cascade properties.," Although the situation we focus on has no external magnetic field, at small scales, the turbulence is still shown as the cascade properties."157 Sell-excited. Alfvén turbulence has also been [ound. (Sokolovetal.2006)., Self-excited $\acute{e}$ n turbulence has also been found \citep{sokolov06}.158. Moreover. we note that the and magnetized turbulence have a high-degree similarity (Cho.Lazarian&2002:Lazarian&Deresnvak 2004).," Moreover, we note that the non-magnetized and magnetized turbulence have a high-degree similarity \citep{cho02,lazarian04}."159. All these evidences indicate that a general form of fluid turbulence can also be valid for the study of random magnetic field generation., All these evidences indicate that a general form of fluid turbulence can also be valid for the study of random magnetic field generation.160" Furthermore. the index of the turbulence spectrum is not universal,"," Furthermore, the index of the turbulence spectrum is not universal."161 Zhou&Matthaeus investigated local turbulent effects with transport models and other nonlinear terms., \citet{zhou90} investigated local turbulent effects with transport models and other nonlinear terms.162 Using (he scaling model (She&Leveque1994). which presents the cascade as an infinitelv-divisible log-Poisson process (She&Wavimire1995).. Doldvrev.Nordlund&Pacloan(2002) derived a steeper spectrum compared to that of IXolmogorov.," Using the scaling model \citep{she94} which presents the cascade as an infinitely-divisible log-Poisson process \citep{she95}, \citet{boldyrev02} derived a steeper spectrum compared to that of Kolmogorov."163 In fact. as estimated by Wane(2002).. (he index value of a turbulent spectrum has the range between —1 to —2.," In fact, as estimated by \citet{wang02}, the index value of a turbulent spectrum has the range between $-1$ to $-2$ ."164 found that the local turbulent spectrum does not show a straight power-law., \citet{maclow00} found that the local turbulent spectrum does not show a straight power-law.165 Therefore. we propose that the turbulent spectrum be shown as:," Therefore, we propose that the turbulent spectrum be shown as:"166svslenms evolve in such way thal we cannot produce models in agreement with the observed values. since alvpical configurations appear.,"systems evolve in such way that we cannot produce models in agreement with the observed values, since atypical configurations appear."167 The mixing in the phase space. lor example. shows us (hat our DM's models do not recover the observational characteristics and. therefore. this model cannot be considered as a good representation of disk galaxies.," The mixing in the phase space, for example, shows us that our BM's models do not recover the observational characteristics and, therefore, this model cannot be considered as a good representation of disk galaxies."168 Bul. one could argue that the snapshot at time /=1 Gyr resembles a typical late-tvpe system. and this model could be considered reliable. since alter /=0.9 ανν this configuration seems to be stable.," But, one could argue that the snapshot at time $t=1$ Gyr resembles a typical late-type system, and this model could be considered reliable, since after $t \gtrsim 0.9$ Gyr this configuration seems to be stable."169 Later on we consider (his issue again by studying the density profile of the clisk., Later on we consider this issue again by studying the density profile of the disk.170 We now compare the rotation curve (hat comes oul [rom our simulations with the one that comes from the centrifugal equilibrium. ie. e707)=rar). where atr) is given by Equation (4)).," We now compare the rotation curve that comes out from our simulations with the one that comes from the centrifugal equilibrium, i.e. $v^2(r)=r\, a(r)$, where $a(r)$ is given by Equation \ref{moffat_aceleration2}) )."171 This rotation curve reads where M(r) is obtained Irom our simulation., This rotation curve reads where $M(r)$ is obtained from our simulation.172 Note that the above equation is just the one used by DM., Note that the above equation is just the one used by BM.173 In Figure 10.. we show the result of this comparison.," In Figure \ref{fig10}, we show the result of this comparison."174 Note the similarity between the rotation curves. even when considered different simulated times.," Note the similarity between the rotation curves, even when considered different simulated times."175 The differences come from parücle noise. verv common in ;/N-body samplings after evolving the system.," The differences come from particle noise, very common in $N$ -body samplings after evolving the system."176 This result shows the consistency between our simulations and the calculation done by DM., This result shows the consistency between our simulations and the calculation done by BM.177 An important issue is (o examine the radial density profile of the simulated disks., An important issue is to examine the radial density profile of the simulated disks.178 A reliable model of a spiral galaxy must have al the end of the simulation (he same initial profile., A reliable model of a spiral galaxy must have at the end of the simulation the same initial profile.179and the ius velocities vary. they all eive the same answer: deusity slope iu the ceutral LO kpc should bea>2.0 and the velocity anisotropy at large distances should be 32 0.7.,"and the rms velocities vary, they all give the same answer: density slope in the central 40 kpc should be $\alpha >-2.0$ and the velocity anisotropy at large distances should be $\beta>0.7$ ."180" Figure 1. shows (dashed curves) the results for the same set of piriuueters as in Angusetal (2007): a= 0. aq =3.lor,=d0 Ἰηρο ry,=(20.10) kpe. oy=1 and 4=2.57."," Figure \ref{fig:Tuned} shows (dashed curves) the results for the same set of parameters as in \citet{Angus2007}: : $\alpha=0$ $\alpha_1=3.1$ , $r_a=40$ kpc, $r_b=(20-40)$ kpc, $\beta_0 =1$ and $\beta_1=-2.57$."181 Tere we use the approximation even in eq.(7))., Here we use the approximation given in \ref{eq:betb}) ).182 The solution is very coutrivect: relatively small deviations from the best hehavior (siuall slope iu the ceuter and radial orbits in the outer radii) result iu failed fits., The solution is very contrived: relatively small deviations from the best behavior (small slope in the center and radial orbits in the outer radii) result in failed fits.183 The full curves in the Figure show what happens when the velocity anisotropy ects less racial: jy=0.5., The full curves in the Figure show what happens when the velocity anisotropy gets less radial: $\beta_0 =0.5$.184 Even this fine-tuned solution fails unless the niass-to-lielit ratio for the larger magnitudebin is arificially increased by a factor of two: 7.5. which eives M.=34100A.," Even this fine-tuned solution fails unless the mass-to-light ratio for the larger magnitudebin is arificially increased by a factor of two: $M/L_g=7.5$ , which gives $M_*=3\times 10^{11}M_{\odot}$."185 Angusetal.(2007). also. found the same trend. but thev made two mistakes. which did not allow them to clearly see the problem.," \citet{Angus2007} also found the same trend, but they made two mistakes, which did not allow them to clearly see the problem."186 First. the solar mass-to-light ratio was used for the D band instead of the ο band.," First, the solar mass-to-light ratio was used for the B band instead of the g band."187 Second. the width of the magnitude bin is presented as an crror in AL/L eiving impression of very large uncertainties in M/L. This is not correct: the statistical unucertaiuntv of the average muinosity of galaxies iu cach bin is very low aud can be neglected.," Second, the width of the magnitude bin is presented as an error in M/L giving impression of very large uncertainties in M/L. This is not correct: the statistical uncertainty of the average luminosity of galaxies in each bin is very low and can be neglected."188" It should be noted that there is rothing special about the galaxies. which are used rere,"," It should be noted that there is nothing special about the galaxies, which are used here."189 The average hnuuinosities differ bv a factor 1.7., The average luminosities differ by a factor 1.7.190 So. it ds not a large difference.," So, it is not a large difference."191 Colors of the oealaxies are practically the same. which then gives je παλιο ML if sve use stellar population models.," Colors of the galaxies are practically the same, which then gives the same M/L if we use stellar population models."192 The rums velocitics are also perfectly consisteut with simple scaling., The rms velocities are also perfectly consistent with simple scaling.193 For example. the ratio of rimis velocitics at the same projected distance of >]| kpe is 1.3 nuplving simple scaling Lox67R.," For example, the ratio of rms velocities at the same projected distance of 70 kpc is 1.3 implying simple scaling $L\propto \sigma^2R$."194 Roughly speaking. we double the luninosity aud vat doubles the stellar mass.," Roughly speaking, we double the luminosity and that doubles the stellar mass."195 This does not work or MIOND: it needs twice 1inore stellar luas., This does not work for MOND: it needs twice more stellar mass.196 This is definitely a problem because there is no justification why ealaxics with the same colors. with the same old population and practically he same DhDuuimositv should have απαλάσαν differcut IMIF.," This is definitely a problem because there is no justification why galaxies with the same colors, with the same old population and practically the same luminosity should have dramatically different IMF."197 The differeuces are very large: most of the stellar mass in the more Iuninous bin should 6 locked up in dwarts with ~0.237... while there is relatively little of those im the lower biu. which Is cousisten with the Kroupa IME.," The differences are very large: most of the stellar mass in the more luminous bin should be locked up in dwarfs with $\sim1980.2M_{\odot}$, while there is relatively little of those in the lower bin, which is consistent with the Kroupa IMF."199 Iu order o luae the argunienut even more clear. we make analysis of velocities in a different way.," In order to make the argument even more clear, we make analysis of velocities in a different way."200 This time we split the sample by stellar mass. but we still keep only red primaries with «or>2.3.," This time we split the sample by stellar mass, but we still keep only red primaries with $u-r >2.3$."201 We make the analysis twice: for satellites L times and for satellites 10 times less bright than the primary galaxy., We make the analysis twice: for satellites 4 times and for satellites 10 times less bright than the primary galaxy.202 There is no systematic difference between the two isolation conditions: within lo the results are the same., There is no systematic difference between the two isolation conditions: within $1\sigma$ the results are the same.203 Figure 5r illustrates this poiut., Figure \ref{fig:Stmass} illustrates this point.204 We select primary galaxies to have the stellar mass in the range M.=(1.6.3.2)10°AZ., We select primary galaxies to have the stellar mass in the range $M_*=(1.6-3.2)10^{11}M\odot$.205 The average stellar iiass is AL.)=2.2«1077AJ. and the average luminosity is (L5=5.1«10197...," The average stellar mass is $\langle M_*\rangle =2.2\times 10^{11}M\odot$, and the average luminosity is $\langle L_g\rangle =5.4\times 10^{10}l_\odot$."206 There is a lint tha more isolated primaries lave slightlyJünger velocities of satellites., There is a hint that more isolated primaries have slightly velocities of satellites.207 Still. the differences are not statistically siguificaut: for radii larger8 than 50 kpe the clifferenees are sinaller than 10 knmi/s. Then we take the observed stellar mass and use it for MOND models aud apply the best tuned paramcters.," Still, the differences are not statistically significant: for radii larger than 50 kpc the differences are smaller than 10 km/s. Then we take the observed stellar mass and use it for MOND models and apply the best tuned parameters."208" We find that parameters sugeested by Angusetal.(2007).fr,=LO kpc.jy= 1.4— 2.57) nuprove the fits as compared with a constaut 2 models."," We find that parameters suggested by \citet{Angus2007} $r_b=40$ kpc,$\beta_0=1$ , $\beta_1=-2.57$ ) improve the fits as compared with a constant $\beta$ models."209 Still.they are not acceptable.," Still,they are not acceptable."210 For example. at 7=150 kpe the," For example, at $R=150$ kpc the"211"(Fisba Optik, Switzerland).","(Fisba Optik, Switzerland)."212 The exit pupil of the first fibre system fills the όθ-μπι core of the second fibre exactly., The exit pupil of the first fibre system fills the $\mu$ m core of the second fibre exactly.213" The scrambling gain (SG) is defined as the ratio between the displacement of the star in front of the input fibre end and the shift of the point spread function (PSF) on the spectrograph detector: where D is the fibre diameter, d the star shift, f the shift of the PSE and F the full width at half maximum of the PSF."," The scrambling gain (SG) is defined as the ratio between the displacement of the star in front of the input fibre end and the shift of the point spread function (PSF) on the spectrograph detector: where $D$ is the fibre diameter, $d$ the star shift, $f$ the shift of the PSF, and $F$ the full width at half maximum of the PSF."214 We measured a value of 1150 for the near field contribution of the scrambler to the SG., We measured a value of 1150 for the near field contribution of the scrambler to the SG.215" For comparison, the SG of the same 60-um fibre without scrambler was measured to be only 123."," For comparison, the SG of the same $\mu$ m fibre without scrambler was measured to be only 123."216 The improved illumination stability of the grating by the scrambler comes at the cost of lower throughput., The improved illumination stability of the grating by the scrambler comes at the cost of lower throughput.217" The scrambler is therefore only implemented on the low-resolution fibre, optimised for high-stability radial velocity work."," The scrambler is therefore only implemented on the low-resolution fibre, optimised for high-stability radial velocity work."218 The sliced high-resolution fibre focuses on high efficiency and does not incorporate a scrambler., The sliced high-resolution fibre focuses on high efficiency and does not incorporate a scrambler.219" We measured the efficiency of the scrambled fibre (LRF) with respect to the HRF to lie around70%;; however, this is not entirely due to the scrambler alone."," We measured the efficiency of the scrambled fibre (LRF) with respect to the HRF to lie around; however, this is not entirely due to the scrambler alone."220 The smaller throughput can be partly explained by the smaller ΤΕΕ fibre diameter (60 versus um)., The smaller throughput can be partly explained by the smaller LRF fibre diameter (60 versus $\mu$ m).221" The telescope interface, installed at the Nasmyth A focal station of the Mercator telescope, links the telescope via the optical fibres to the spectrograph."," The telescope interface, installed at the Nasmyth A focal station of the Mercator telescope, links the telescope via the optical fibres to the spectrograph."222" It includes an atmospheric dispersion corrector (ADC), the fibre entrance, a fibre viewer/telescope guiding system, the calibration light projection optics, and a mask that selectively covers one or both of the star fibres (Fig. 11))."," It includes an atmospheric dispersion corrector (ADC), the fibre entrance, a fibre viewer/telescope guiding system, the calibration light projection optics, and a mask that selectively covers one or both of the star fibres (Fig. \ref{fig:telescope_interface}) )."223" Although the fibre entrances have a fairly large sky aperture of 2.15 or aarcsec, HERMES would be limited to observations at zenithal angles smaller than 50? if we want to avoid losing an important part of the far blue or red flux due to differential atmospheric dispersion."," Although the fibre entrances have a fairly large sky aperture of 2.15 or arcsec, HERMES would be limited to observations at zenithal angles smaller than $50^{\circ}$ if we want to avoid losing an important part of the far blue or red flux due to differential atmospheric dispersion."224 The HERMES atmospheric dispersion corrector provides a step-wise correction of the atmospheric dispersion by means of a set of four correctors in a selection wheel., The HERMES atmospheric dispersion corrector provides a step-wise correction of the atmospheric dispersion by means of a set of four correctors in a selection wheel.225" Each corrector consists of a cemented doublet prism of N-BK7 and ΕΙ glass (wedge angles: 0.55°, 1.85?, 3.60? and 5.40°)."," Each corrector consists of a cemented doublet prism of N-BK7 and LLF1 glass (wedge angles: $0.55^{\circ}$, $1.85^{\circ}$, $3.60^{\circ}$ and $5.40^{\circ}$ )."226 The corrector selection matches the requirement that the length of the secondary spectrum never becomes larger than aarcsec for zenithal angles up to 65° and that it stays below 2aarcsec up to 72°., The corrector selection matches the requirement that the length of the secondary spectrum never becomes larger than arcsec for zenithal angles up to $65^{\circ}$ and that it stays below arcsec up to $72^{\circ}$.227" The fibres, together with their micro lenses, are precisely mounted at the back of a polished stainless steel mirror in the focal plane of the telescope, just in front of the 175-um and 150-um diameter apertures (Fig."," The fibres, together with their micro lenses, are precisely mounted at the back of a polished stainless steel mirror in the focal plane of the telescope, just in front of the $\mu$ m and $\mu$ m diameter apertures (Fig."228 10 a)., \ref{fig:fibre_entrance} a).229 This concave mirror is slightly inclined (8?) to reflect the image of the field via a 45?- mirror through the fibre-viewer optics to the guiding CCD., This concave mirror is slightly inclined $8^{\circ}$ ) to reflect the image of the field via a $45^{\circ}$ -fold mirror through the fibre-viewer optics to the guiding CCD.230" The fibre-viewer optics reduce the telescope f/12 focal ratio to f/6.5, providing a comfortable acquisition field of 66 arcmin on an SBIG ST-1603ME CCD camera 11530 9-um pixels)."," The fibre-viewer optics reduce the telescope $f/12$ focal ratio to $f/6.5$, providing a comfortable acquisition field of 6 arcmin on an SBIG ST-1603ME CCD camera 1530 $\mu$ m pixels)."231" During acquisition, the image of the star is centred on one star-fibre hole, while a movable mask blacks out the other fibre hole."," During acquisition, the image of the star is centred on one star-fibre hole, while a movable mask blacks out the other fibre hole."232 The light of the wings of the PSF that will not enter the star-fibre hole is subsequently used for telescope guiding., The light of the wings of the PSF that will not enter the star-fibre hole is subsequently used for telescope guiding.233 A separate calibration unit is connected to the telescope interface through a 300-ym optical fibre., A separate calibration unit is connected to the telescope interface through a $\mu$ m optical fibre.234" The calibration light projection optics, mounted on a linear translation stage, can inject calibration light into each of the star fibres separately, thus providing optimal flat-field and wavelength calibration possibilities."," The calibration light projection optics, mounted on a linear translation stage, can inject calibration light into each of the star fibres separately, thus providing optimal flat-field and wavelength calibration possibilities."235 The projection optics are designed to deliver a, The projection optics are designed to deliver a236spiral sample. along with a similar number of the random mass niatched blue spirals.,"spiral sample, along with a similar number of the random mass matched blue spirals."237 We find that red spirals have an optical bar fraction of at least 6725% (raw fractions were T4 lor BC and 67% for WLAL) and up to including more uncertain identifications. while the bluc spirals have a xw fraction of only 27£5% [rom BC and from IXLM).," We find that red spirals have an optical bar fraction of at least $67\pm5\%$ (raw fractions were $72\%$ for BC and $67\%$ for KLM) and up to including more uncertain identifications, while the blue spirals have a bar fraction of only $27\pm5\%$ from BC and from KLM)."238 These bar identifications used the SDSS gri images vpically used. by. Galaxy Zoo. and. were based. on classic visual bar identification methods such as those used. by deVaucouleursetal.(L991) to find a bar fraction of in he RCS.," These bar identifications used the SDSS $gri$ images typically used by Galaxy Zoo, and were based on classic visual bar identification methods such as those used by \citet{RC3} to find a bar fraction of in the RC3."239 More recent work on bar fractions in the literature (c.g.Jogeeοἱal.04:ShethctOS:Darazzaet2008:indingAguerrietal.2009) use automated. techniques. [or bars using elliptical isophote fitting.," More recent work on bar fractions in the literature \citep[e.g.][]{J04,KS08,Bar08,A09} use automated techniques for finding bars using elliptical isophote fitting."240" The two studies using SDSS data on ~2000 ""disk"" galaxies (Barazzactal.2008:Aguerrietal.2009) are most directly comparable o this study."," The two studies using SDSS data on $\sim$ 2000 “disk"" galaxies \citep{Bar08,A09} are most directly comparable to this study."241 Both however use automated. techniques to identifv disk/spiral samples based. largely on concentration (Aguerrietal.2009). and colour (Barazzactal...2008.thisromcolourselected“spirals” onlv).. so we expect svstematic dillerences with our visually identified face-on spiral sample and inparticular point out that red spirals by our definition will be completely absent. from the Barazzaetal.(2008) study. and extremely rare in. Ag@uerrietal.(2009).," Both however use automated techniques to identify disk/spiral samples based largely on concentration \citep{A09} and colour \citep[][this study does also consider Sersic fits, but the final results are from colour selected ``spirals"" only]{Bar08}, so we expect systematic differences with our visually identified face-on spiral sample and inparticular point out that red spirals by our definition will be completely absent from the \citet{Bar08} study, and extremely rare in \citet{A09}."242.. Phe overall fractions found by Barazzaetal.(2008):Agucrrietal.(2009)M are comparable to our total bar fraction. of ~5O% (they find and of their disk samples hosting xus). but intriguinglv both studies suggest a trend For blucr clisk galaxies to be more likely to host bars. in cirect contrast o the clear signal we find for red spirals have more obvious bars in the gr images.," The overall bar fractions found by \citet{Bar08,A09} are comparable to our total bar fraction of $\sim$ (they find and of their disk samples hosting bars), but intriguingly both studies suggest a trend for bluer disk galaxies to be more likely to host bars, in direct contrast to the clear signal we find for red spirals to have more obvious bars in the $gri$ images."243 Lt is not clear at this point if he difference in these trends is due to the sample selection or he bar identification method (although Aguerrietal.(2009) claim only a dilerence between their automated. bar inder and a visual check of their sample)., It is not clear at this point if the difference in these trends is due to the sample selection or the bar identification method (although \citet{A09} claim only a difference between their automated bar finder and a visual check of their sample).244 A more detailed study of bar fractions in Galaxy Zoo galaxies as a function of various galaxy properties and considering possible biases on the visual bar identification method with gri images is being prepared using bar identifications for almost. 30.000 spiral galaxies collected« the second. phase of Galaxy Zoo (GZ2: Masters et. al.," A more detailed study of bar fractions in Galaxy Zoo galaxies as a function of various galaxy properties and considering possible biases on the visual bar identification method with $gri$ images is being prepared using bar identifications for almost 30,000 spiral galaxies collected during the second phase of Galaxy Zoo (GZ2; Masters et al."245 Mep.)ο, in prep.).246 For the purposes of this work. the huge increase in bar fraction. between the blue and red. spiral samples gives such a strong hint of a trend of bar fraction with colour(in the sense that red spirals are much more likely to host bars). that it suggests bars may be providing an important clue to the formation of the red spirals ancl therefore the impact of bars should be considered in a discussion of their possible origin.," For the purposes of this work, the huge increase in bar fraction between the blue and red spiral samples gives such a strong hint of a trend of bar fraction with colour (in the sense that red spirals are much more likely to host bars), that it suggests bars may be providing an important clue to the formation of the red spirals and therefore the impact of bars should be considered in a discussion of their possible origin."247 In simulations of spiral galaxy formation. bars [form quickly once a stable disk. is Formed. and are dillicult to clestrov(ο.Debattistaetal.2006).," In simulations of spiral galaxy formation, bars form quickly once a stable disk is formed, and are difficult to destroy \citep[e.g.][]{D06}."248 Mowever. the impact of higher density environments on bars is unclear.," However, the impact of higher density environments on bars is unclear."249 Vidal night induce bar formation (eg., Tidal interactions might induce bar formation (eg.250" Hernaquist AMihos MN199. but they also act to heat the disks of spirals. and bars "" most quicklyin cold disks(eg."," Hernquist Mihos 1995), but they also act to heat the disks of spirals, and bars form most quickly in cold disks (eg."251 Toomre 1964)., Toomre 1964).252 If external triggers such as tidal interactions are the most important source of bar instabilities. then the higher density environments of the red spirals may naturally lead to high bar fractions.," If external triggers such as tidal interactions are the most important source of bar instabilities, then the higher density environments of the red spirals may naturally lead to high bar fractions."253 A possible explanation for the high bar fraction in red spirals could be that the bars themselves are at least: partly responsible for the process which turned olf star formation in these objects., A possible explanation for the high bar fraction in red spirals could be that the bars themselves are at least partly responsible for the process which turned off star formation in these objects.254 Bars are known to be the most ellicient way to redistribute material in the disks of galaxies (eg., Bars are known to be the most efficient way to redistribute material in the disks of galaxies (eg.255 Combes Sanders 1981). by channelling eas into the central regions of the galaxy.," Combes Sanders 1981), by channelling gas into the central regions of the galaxy."256 Bars have been invoked. (eg., Bars have been invoked (eg.257 Shlosman et n, Shlosman et al.258 2000) as à wav to feed. gas to the central black hole -n the increase in LINEN. fraction we observe in the red s could be a remnant of this process (if they are LINERs associated with ACN)., 2000) as a way to feed gas to the central black hole - the increase in LINER fraction we observe in the red spirals could be a remnant of this process (if they are LINERs associated with AGN).259 Perhaps the bars in the red. spirals mave removed the cold gas from the disk and channelled it inwards where it has either been used as fuel for the AGN or created a starburst.," Perhaps the bars in the red spirals have removed the cold gas from the disk and channelled it inwards where it has either been used as fuel for the AGN, or created a starburst."260 Of course. this starburst must have lappencc more than I yr ago. to be consistent with our observation that red. spirals are not post starburst objects. out if bars are as robust as simulations suggest then they would. still persist for a long time after evidence of any bar rigecred central starburst was removed.," Of course, this starburst must have happened more than $\sim1$ Gyr ago, to be consistent with our observation that red spirals are not post starburst objects, but if bars are as robust as simulations suggest then they would still persist for a long time after evidence of any bar triggered central starburst was removed."261 Since red spirals are observed at. all density levels. the process which creates red spirals cannot. be confined. only to regions of high galaxy density: environment. alone is not sullicient to determine whether a galaxy will become a red spiral or not.," Since red spirals are observed at all density levels, the process which creates red spirals cannot be confined only to regions of high galaxy density; environment alone is not sufficient to determine whether a galaxy will become a red spiral or not."262 The lack of any clear correlation of the star formation rates of red spirals with environment. also indicates that τος spirals in low clensitw regions are similar to those in clusters making it more dillicult to invoke only environmental processes in their formation.," The lack of any clear correlation of the star formation rates of red spirals with environment, also indicates that red spirals in low density regions are similar to those in clusters making it more difficult to invoke only environmental processes in their formation."263 The process which turns spirals red may be more likely to occur in higher density regions. but. must. be. possible.way dn alb environments. unless completely different mechanisms are responsible for red spirals in high and low density regions.," The process which turns spirals red may be more likely to occur in higher density regions, but must be possible, in all environments, unless completely different mechanisms are responsible for red spirals in high and low density regions."264 That the red. spiral population has a significantly higher mean mass than the population of blue spirals. also indicates that they are not. uniformly sourced. from the field. population.," That the red spiral population has a significantly higher mean mass than the population of blue spirals, also indicates that they are not uniformly sourced from the field population."265 However. perhaps the mass (transition between red and blue spirals represents the lowest mass spiral which can retain its spiral structure uncer the inlluence of environmental cllects which shut off its star formation," However, perhaps the mass transition between red and blue spirals represents the lowest mass spiral which can retain its spiral structure under the influence of environmental effects which shut off its star formation."266 low mass blue spirals might. pass through the red spiral phase very quickly. or experience both morphological and colour transformation at the same time).," low mass blue spirals might pass through the red spiral phase very quickly, or experience both morphological and colour transformation at the same time)."267" We sugeest that rather than representing an intermediate stage of environmental transformation. red. spirals could. be ""old. spirals” who through normal internal evolution have already used up their reservoirs of gas - perhaps aided by the redistribution of gas due to a xw instabilitv."," We suggest that rather than representing an intermediate stage of environmental transformation, red spirals could be “old spirals"" who through normal internal evolution have already used up their reservoirs of gas - perhaps aided by the redistribution of gas due to a bar instability."268. We suggest. that part of the reason they are found. to be more common in higher density. regions is because the initial density Uuctuations decoupled. from he Hubble Gow earlier there (seeBrommetal.2009.forareview) anc galaxies started. assembling at much earlier ines (Cooperctal.2009). in dense environments so have iac longer to use up their gas., We suggest that part of the reason they are found to be more common in higher density regions is because the initial density fluctuations decoupled from the Hubble flow earlier there \citep[see][for a review]{VB09} and galaxies started assembling at much earlier times \citep{C09} in dense environments so have had longer to use up their gas.269 Such objects would naturally x found at the high mass end of the galaxy. clistribution as they have been assembling for a long time., Such objects would naturally be found at the high mass end of the galaxy distribution as they have been assembling for a long time.270 Rec spirals oesumablv become less common in the highest. densities where strong environmental processes are more important and may have already changed all spirals into early tvpe ealaxies., Red spirals presumably become less common in the highest densities where strong environmental processes are more important and may have already changed all spirals into early type galaxies.271The UV coutiuua and spectral features of 1531. IRÁASITIGU-3111. (SAO 209306) aud. IRASIS379-1:07 (LSS 5112) were in good agreement with the dereddened UV(GIUE) spectra of standard stars (Table 3) of simular optica spectral types (Fig.,"The UV continua and spectral features of IRAS17203-1534, IRAS17460-3114 (SAO 209306) and IRAS18379-1707 (LSS 5112) were in good agreement with the dereddened UV(IUE) spectra of standard stars (Table 3) of similar optical spectral types (Fig."272 2)., 2).273 The V) values of hese stars determined from the feature are nearly the same as Vous (Table 3) sugeestiug ucelieible extinction of starlight duc to ciremustellar dust in these three cases., The $-$ V) values of these stars determined from the feature are nearly the same as $-$ $_{\rm total}$ (Table 3) suggesting negligible extinction of starlight due to circumstellar dust in these three cases.274 Emission lues of Sill(1533. 1505. ). ). } and Fell(17sh. ) in the spectrum of IRASI258I-1837 (ITen23-817) indicate the presence of hot plasma or a nebula.," Emission lines of SiII(1533, 1808, ), ), ) and FeII(1785, ) in the spectrum of IRAS12584-4837 (Hen3-847) indicate the presence of hot plasma or a nebula."275 The UV(IUE) spectra of LO stars (IRASI3266-5551 (CPD-55 Shas). IRASII331-6135. (IIen3-1013).. IRÀSI6202-595G (SAO 213756). IRASIT07L-1815.— (Hou3-1317). IRASI17311-1921. (Ilen2-1128).. IRASISO23-3109. (LSS 1631)... IRASISO0G2|2110. (SAO s5766). IRASIS3TI-3159 (LSE ο). IRAS22023|5219 (LSIII 15221). TRAS22195|5131 (LST 15112)). dereddened using the feature in the UV. showed considerably reddened. continua dui comparison with the dereddened spectra of standard stars of simular optical spectral tvpes (Figs.," The UV(IUE) spectra of 10 stars (IRAS13266-5551 (CPD-55 5588), IRAS14331-6435 (Hen3-1013), IRAS16202-5956 (SAO 243756), IRAS17074-1845 (Hen3-1347), IRAS17311-4924 (Hen3-1428), IRAS18023-3409 (LSS 4634), IRAS18062+2410 (SAO 85766), IRAS18371-3159 (LSE 63), IRAS22023+5249 (LSIII +5224), IRAS22495+5134 (LSIII +5142)), dereddened using the feature in the UV, showed considerably reddened continua in comparison with the dereddened spectra of standard stars of similar optical spectral types (Figs."276 3.5 and 6).," 3, 5 and 6)."277 Comparing the interstellar extinction estimates from the feature with the total extinction νι). towards these stars. we find that these stars have cousicderable UV deficiency and circtuustcellar extinction.," Comparing the interstellar extinction estimates from the feature with the total extinction $-$ $_{\rm total}$ ) towards these stars, we find that these stars have considerable UV deficiency and circumstellar extinction."278 The hot ceutral star of the bipolar proto-planctary nebula (PPN). TRASL7123-1755 (Πο175) was not detected iu a 35 müuute exposure with the SWP camera.," The hot central star of the bipolar proto-planetary nebula (PPN), IRAS17423-1755 (Hen3-1475) was not detected in a 35 minute exposure with the SWP camera."279 This may be due to obscuration of the ceutral star by a dusty disk., This may be due to obscuration of the central star by a dusty disk.280 UST WFEDPC? images of the object showed the presence of a dusty torus with a spatial exteut of 2” (Borkowski ct al..," HST WFPC2 images of the object showed the presence of a dusty torus with a spatial extent of $\arcsec$ (Borkowski et al.,"281 1997)., 1997).282 To account for the observed UW deficiency in the 10 hot post-AGB candidates mentioned in Sec., To account for the observed UV deficiency in the 10 hot post-AGB candidates mentioned in Sec.283 3.3 and to understaud the shape of the UV σοι im these stars. we investigated the circumstellar extinction law in these cases.," 3.3 and to understand the shape of the UV continuum in these stars, we investigated the circumstellar extinction law in these cases."284 Waters et al. (, Waters et al. (2851989) modelled. the circumstellar extinction in the case of the post-ACB star. IIR1019.,"1989) modelled the circumstellar extinction in the case of the post-AGB star, HR4049."286 We followed he sane procedure here., We followed the same procedure here.287" We plotted the logarithmic difference between the dereddened UW flux. of a hot post-ACD candidate (normalised to its V-band flux) aud the dereddened UV flux of the corresponding standard star (normalised to its V-banud flux). A=log(tyο...οf-)randard Vs. Af, "," We plotted the logarithmic difference between the dereddened UV flux of a hot post-AGB candidate (normalised to its V-band flux) and the dereddened UV flux of the corresponding standard star (normalised to its V-band flux), $\Delta = {\rm log (f_{\lambda}/f_{v})_{star} - log (f_{\lambda}/f_{v})_{standard}}$ Vs. $\lambda ^{-1}$."288Fie., Fig.289 I shows the plot of the logarithmic flux deficiency due to ciretuustellar dust from 3.2 to 8 p.1 for the 10 stars., 4 shows the plot of the logarithmic flux deficiency due to circumstellar dust from 3.2 to 8 $\mu^{-1}$ for the 10 stars.290 Best fit lines were obtained by minimising the chi-square error statistic., Best fit lines were obtained by minimising the chi-square error statistic.291 Ee..," Eg.,"292" in the case of IRASLT311-1921 (Hen3-1128). we obtained. A = 0.07 O.L0A2,"," in the case of IRAS17311-4924 (Hen3-1428), we obtained, $\Delta$ = 0.07 $-$ $\lambda^{-1}$."293 To derive the ciretuustellar extinction Vocus.) dn magnitudes (Tables la aud b) A had to be multiplied by 2.5 and A = Lge was used., To derive the circumstellar extinction $-$ $_{\rm C.S.}$ ) in magnitudes (Tables 4a and b) $\Delta$ had to be multiplied by $-$ 2.5 and $\lambda$ = $\mu$ was used.294 For TRAS$22023|529 (LSIII |5221). only a D spectral type 1s isted iu Literature.," For IRAS22023+5249 (LSIII +5224), only a B spectral type is listed in literature."295 We compared the spectrum of this star with that of a B2-supereiant standard star., We compared the spectrum of this star with that of a B2-supergiant standard star.296 TRAS2?2195|5134 (LSIIT 15112) was detected as a PN with an angular extent of 0.57(Tyleuda Stasiuska. 1991).," IRAS22495+5134 (LSIII +5142) was detected as a PN with an angular extent of (Tylenda Stasinska, 1994)."297 Central stars of PNe have temperatures n excess of ~ 00001 corresponding to spectral types of O9 or hotter., Central stars of PNe have temperatures in excess of $\sim$ 30000K corresponding to spectral types of O9 or hotter.298 We compared the UV(IUE) spectra of this star with a standard OOV star (IID38666)., We compared the UV(IUE) spectra of this star with a standard O9V star (HD38666).299 For each of the 10 hot post-ACB candidates we found that the ciremustellar extinction varies as \1 (Fig., For each of the 10 hot post-AGB candidates we found that the circumstellar extinction varies as $\lambda^{-1}$ (Fig.300 1)., 4).301 The derived Viers. values in Table la account well for the difference between the total and the interstellar extinction (from the feature) values., The derived $-$ $_{\rm C.S.}$ values in Table 4a account well for the difference between the total and the interstellar extinction (from the feature) values.302 Vocus. values in Table tb are in excess of the difference between Vs and the interstellar extinction from, $-$ $_{\rm C.S.}$ values in Table 4b are in excess of the difference between $-$ $_{\rm total}$ and the interstellar extinction from.303 This may be because of the variable nature of these starpA and because the V and Vj magnitudes at cach epoc[um of the IUE observations are not kuown., This may be because of the variable nature of these stars and because the V and $-$ V) magnitudes at each epoch of the IUE observations are not known.304 Mean V aud V)maguitudes frou literature have been used for cach of these two stars (see Sec., Mean V and $-$ V) magnitudes from literature have been used for each of these two stars (see Sec.305 3.3.2 below)., 3.3.2 below).306 IRAS 16206-5956 was found to be variable in the UV (Fig., IRAS 16206-5956 was found to be variable in the UV (Fig.307 1)., 1).308 The spectrum of the star has chaueed from 21 July 1988 to 12 March 1991 and from 12 March 1991 to 28 April 1991 suggesting both lone term and short term variability., The spectrum of the star has changed from 21 July 1988 to 12 March 1994 and from 12 March 1994 to 28 April 1994 suggesting both long term and short term variability.309 The παπα flux in the UV was observed ou 12 March 1991., The maximum flux in the UV was observed on 12 March 1994.310 Schild et al (1983) found it to be variable in the optical with AV=0.13., Schild et al (1983) found it to be variable in the optical with $\Delta$ V=0.13.311 Since the V imagnitudes of the star at the epochs of the UV(IUE) obscrvatious are uot known. we used aiean V maeuitude (=9.76) from the photometric and spectroscopic database for Steplieusonu-Sanduleak Lunuinous Stars in the Southern Milkv. Wav (ποσα. 1998).," Since the V magnitudes of the star at the epochs of the UV(IUE) observations are not known, we used a mean V magnitude (=9.76) from the photometric and spectroscopic database for Stephenson-Sanduleak Luminous Stars in the Southern Milky Way (Reed, 1998)."312 Its spectral type iu literature is listed as A3labe (Thunphrevs. 1975. Parthasarathy et aL.," Its spectral type in literature is listed as A3Iabe (Humphreys, 1975, Parthasarathy et al.,"313 2000a) and AQIae (Schild et al., 2000a) and A0Iae (Schild et al.314 1983. Carrisou ct al.," 1983, Garrison et al.,"315 1971)., 1977).316 Oucunaijer (1996) listed it as Bala. citing the Simbad database.," Oudmaijer (1996) listed it as B8Ia, citing the Simbad database."317 However. we could not find a reference for the same.," However, we could not find a reference for the same."318 We compared the dereddeued IUE spectra of the star at different epochs with the derccddened spectra of ΑΟ (ΠΟΙΟ025) and AOTa (ITD21289). standard stars, We compared the dereddened IUE spectra of the star at different epochs with the dereddened spectra of A3Ib (HD104035) and A0Ia (HD21389) standard stars319In this section. we look at the impact of photometric redshift estimation on dark energy science.,"In this section, we look at the impact of photometric redshift estimation on dark energy science."320 In. particular. we concentrate on the measurement of dark energy. using galaxy power spectra and barvon acoustic oscillations.," In particular, we concentrate on the measurement of dark energy using galaxy power spectra and baryon acoustic oscillations."321 The galaxy angular power spectrum: is à measure of he clustering in the galaxy population within time bins extending from the present to a time when the Universe was only a third of its present age., The galaxy angular power spectrum is a measure of the clustering in the galaxy population within time bins extending from the present to a time when the Universe was only a third of its present age.322 Large-scale surveys like DES »ovide ideal data sets for studying the elustering properties of galaxies and therefore the clustering properties of their uncderlving dark matter distribution and hence are useful xobes for mapping how the dark matter distribution evolves with time., Large-scale surveys like DES provide ideal data sets for studying the clustering properties of galaxies and therefore the clustering properties of their underlying dark matter distribution and hence are useful probes for mapping how the dark matter distribution evolves with time.323 Furthermore. many other characteristic features appear in the power spectrum which provide hat can be used to determine the angular cliameter distance. Dy. as a function of redshift.," Furthermore, many other characteristic features appear in the power spectrum which provide that can be used to determine the angular diameter distance, $D_A$, as a function of redshift."324 Barvon acoustic oscillations are one such feature of interest which appear as on the rower spectrum., Baryon acoustic oscillations are one such feature of interest which appear as on the power spectrum.325 The position of the peaks anc troughs of hese wigeles in Fourier space can be used to determine a set of cosmological parameters e.g. ?. and ?.., The position of the peaks and troughs of these wiggles in Fourier space can be used to determine a set of cosmological parameters e.g. \citet{Blake:03} and \citet{Seo:03}.326 The accuracy with which we can measure this tvpical acoustic scale is proportional to the average fractional error in the power spectrum. 02/2.," The accuracy with which we can measure this typical acoustic scale is proportional to the average fractional error in the power spectrum, $\delta P/P$."327 Lhe fractional error on the power spectrum arises from two sources., The fractional error on the power spectrum arises from two sources.328 Firstly. the number of independant spatial modes that we can measure in a given volume is finite and this will lead to errors in the power spectrum that are proportional to L/W.," Firstly, the number of independant spatial modes that we can measure in a given volume is finite and this will lead to errors in the power spectrum that are proportional to $1/\sqrt{V}$."329 This is known as cosmic variance., This is known as cosmic variance.330 Secondly. there is a contribution from shot noise due to imperfect sampling of the fluctuations as we only have a finite number of tracers of these fluctuations within a given volume.," Secondly, there is a contribution from shot noise due to imperfect sampling of the fluctuations as we only have a finite number of tracers of these fluctuations within a given volume."331 Lo we assume a density. field. that follows Ciaussian statistics. we can follow ? and assume theerror on the power spectrum measurement. £2? is weighted in the following wav: where m is the mean number density in a given volume as seen by an observer and can be written in terms of the galaxy redshift distribution as follows: where fai is the fraction of the sky. covered by the survey and dVdz is the comoving volume clement.," If we assume a density field that follows Gaussian statistics, we can follow \citet{FKP:PowerSpec} and assume theerror on the power spectrum measurement, $P$ is weighted in the following way: where $n$ is the mean number density in a given volume as seen by an observer and can be written in terms of the galaxy redshift distribution as follows: where $f_{sky}$ is the fraction of the sky covered by the survey and $dV/dz$ is the comoving volume element."332 The first term in Eq., The first term in Eq.333 7 denotes the οσοι of cosmic variance while the second term is the contribution [rom shot noise., \ref{eq:errorps} denotes the effect of cosmic variance while the second term is the contribution from shot noise.334 In order to minimise the error on the power spectrum. one has to design a survey with maximum volume provided there are enough. sources within this volume for the shot noise contribution to be minimal.," In order to minimise the error on the power spectrum, one has to design a survey with maximum volume provided there are enough sources within this volume for the shot noise contribution to be minimal."335 I£ 54?73 the power spectrum is well estimated and there is no significant advantage to be gained with more galaxies (2).., If $nP>3$ the power spectrum is well estimated and there is no significant advantage to be gained with more galaxies \citep{Seo:03}.336 In this work. we assume that to obtain a reasonable estimate of the power spectrum we need to satisfy the condition. nP1.," In this work, we assume that to obtain a reasonable estimate of the power spectrum we need to satisfy the condition, $nP \gtrsim 1$."337" Taking into account the galaxy bias. 6 that scales the galaxy. power spectrum to the matter power spectrum. and including the scaling of the matter power spectrum with redshift as a linear growth factor. D(z) we get the following expression [or nuu We have used the formalism for the transfer function set out in 7. to calculate our power spectrum at A,=0.15AJpe as this is well within the linear regime of the power spectrum."," Taking into account the galaxy bias, $b$ that scales the galaxy power spectrum to the matter power spectrum, and including the scaling of the matter power spectrum with redshift as a linear growth factor, $D(z)$ we get the following expression for $nP_{gal}$: We have used the formalism for the transfer function set out in \citet{EH:transfer} to calculate our power spectrum at $k_*=0.1h{Mpc}^{-1}$ as this is well within the linear regime of the power spectrum."338 At larger values of &. non-linearities due to clustering and other structure formation start to dominate and make it harder to detect the BAO signal.," At larger values of $k$, non-linearities due to clustering and other structure formation start to dominate and make it harder to detect the BAO signal."339" We assume a survey with θ«z2 and for,=0.119.", We assume a survey with $0<z<2$ and $f_{sky}=0.119$.340 The bias is assumed. to be 127., The bias is assumed to be 1.2.341". In Figure 12... we plot nD, asa function of the redshift."," In Figure \ref{fig:PowerSpec}, we plot $nP_{gal}$ as a function of the redshift."342 This is done for the entire catalogue and for clipped catalogues with dillerent. clipping thresholds., This is done for the entire catalogue and for clipped catalogues with different clipping thresholds.343 We perform the same analysis for optical only DIES data as well as optical ancl NUR data from DES|VILIS., We perform the same analysis for optical only DES data as well as optical and NIR data from DES+VHS.344 The results are summarised in Table 5 and Figure 12.., The results are summarised in Table \ref{tab:powerspec} and Figure \ref{fig:PowerSpec}.345 From these results we can see that applving a threshold error at which to cut our photometric redshilt catalogue proves cllective in removing outliers [from our. sample before performing any kind. of cosmological analysis on it., From these results we can see that applying a threshold error at which to cut our photometric redshift catalogue proves effective in removing outliers from our sample before performing any kind of cosmological analysis on it.346 For the DES catalogue of redshifts obtained. using grizY photometry. we can remove all galaxies with a threshold error of more than ~0.03 in order to obtain an accurate measurement of the galaxy. power spectrum out to a redshift of 1.," For the DES catalogue of redshifts obtained using $grizY$ photometry, we can remove all galaxies with a threshold error of more than $\sim0.03$ in order to obtain an accurate measurement of the galaxy power spectrum out to a redshift of 1."347 This leaves us with only of our original sample but this sample has a scatter on its photometric redshift that isa [actor of 2.7 times better than that of the original sample and is therefore more effective in constraining the cosmology., This leaves us with only of our original sample but this sample has a scatter on its photometric redshift that is a factor of $\sim2.7$ times better than that of the original sample and is therefore more effective in constraining the cosmology.348 For the DES |. VIIS catalogue. a threshold. error of 0.025 can be applied. to effectively. constrain the galaxy power spectrum to redshift 1.," For the DES + VHS catalogue, a threshold error of 0.025 can be applied to effectively constrain the galaxy power spectrum to redshift 1."349 This leaves us with only of our original sample. but the overall scatter on the photometric redshift has been reduced. by a [actor of ~2.75.," This leaves us with only of our original sample, but the overall scatter on the photometric redshift has been reduced by a factor of $\sim2.75$."350 This is equivalent to. performing an LRG selection on our survey as these galaxies have more accurate large-scale. structure signals ancl more accurate photometric redshifts due to the prominence of their break (??)..," This is equivalent to performing an LRG selection on our survey as these galaxies have more accurate large-scale structure signals and more accurate photometric redshifts due to the prominence of their break \citep{Blake:CosmoLRG,Pad:LRGphotoz}."351 In order to provide a reasonable measurement of the galaxy power spectrum for the entire DIES redshift range of 0<22. we can apply a threshold error cut of z-0.1 to the DIES only sample ancl use most of the galaxies in our analysis.," In order to provide a reasonable measurement of the galaxy power spectrum for the entire DES redshift range of $0<z<2$, we can apply a threshold error cut of $>0.1$ to the DES only sample and use most of the galaxies in our analysis."352 When we add NER. photometry from Ες to our sample. a less conservative clipping eut of 0.05 can be applied and only of the galaxies used to reduce the scatter on the photometric redshift by a factor of 2.," When we add NIR photometry from VHS to our sample, a less conservative clipping cut of 0.05 can be applied and only of the galaxies used to reduce the scatter on the photometric redshift by a factor of $\sim2$."353 Note that as is also reduced in these cases although not to the same extent as the reduction in ea., Note that $\sigma_{68}$ is also reduced in these cases although not to the same extent as the reduction in $\sigma$ .354 A reduction in mx corresponds to a reduction in the intrinsic scatter of our sample minus the outliers., A reduction in $\sigma_{68}$ corresponds to a reduction in the intrinsic scatter of our sample minus the outliers.355 o clippingelippi ourο cataloguescatal wo]in. thishi wayavo before performing any kind. of cosmological analysis on them.," By clipping our catalogues in this way before performing any kind of cosmological analysis on them,"356"that this assumption is well performed at atuosphere depth f>700 e/em?. primary cuereics E>lo’ GeV and zenith angles 0«359, Let also the measurements aud further evaluations of EAS electron auc truncated imon sizes be carried out without errors (GAG/d/N=6(NΑΙ) and inteeral Eq. (","that this assumption is well performed at atmosphere depth $t>700$ $^2$, primary energies $E>10^5$ GeV and zenith angles $\theta<35^0$ Let also the measurements and further evaluations of EAS electron and truncated muon sizes be carried out without errors $dG/dN^*\equiv\delta(N-N^*)$ ) and integral Eq. ("3571) include only statistical Then a set of Eq. (,1) include only statistical Then a set of Eq. (358"1) traustorms iuto the followine: Let's deteune the pariuueters of distribution functions ΟΠ/0NN,.,, by the following kuown enipirical expressions: where the values of corresponding approximation parameters (0....0.0.... p) are presented iu Tables 1.2 aud obtained by CORSIRAGOIGONIKC) EAS simulation code |15] at QGSJET interaction model [16]..","1) transforms into the following: Let's determine the parameters of distribution functions $\partial359W/\partial N_{e,\mu}$ by the following known empirical expressions: where the values of corresponding approximation parameters $a,\dots d,\alpha,\dots\rho$ ) are presented in Tables 1,2 and obtained by CORSIKA6016(NKG) EAS simulation code \cite{CORSIKA} at QGSJET interaction model \cite{QGSJET}."360 The accuracy of approximation (I) for average EAS electron size «ΑννALO)> is less than 10 at 3-10E3-105 GeV. A=1....56. 0«329 and observation level 1020 &/cau?.," The accuracy of approximation (4) for average EAS electron size $<N_e(E,A,\theta)>$ is less than $\%$ at $3\cdot10^5<E<3\cdot10^8$ GeV, $A\equiv1,\dots56$, $\theta<32^0$ and observation level 1020 $^2$."361" The correspouding aceuracies οσο and average EAS tyuucated umou size |, are less than 1-24. Changing the variables of kerucl fuuctious of Eq. ("," The corresponding accuracies of $\sigma_{e,\mu}$ and average EAS truncated muon size $N_{\mu}$ are less than $\%$ Changing the variables of kernel functions of Eq. ("362"5) from AN, and IN, to ic aud e, respectively according to",3) from $N_e$ and $N_{\mu}$ to $x_e$ and $x_\mu$ respectively according to363We only considered temperatures between IIS. aud IIs. For ETE. we uced sutiicicutly high temperature and sufficiently high deusity so that the vibrational levels can be collisionally excited.,"We only considered temperatures between K and K. For LTE, we need sufficiently high temperature and sufficiently high density so that the vibrational levels can be collisionally excited."364 This requires temperatures Του 220004KI/ and densities ereater than wy>10M 7: but temperatures should. also be less than TIS. otherwise CO molecules would dissociate (?)..," This requires temperatures $_{\rm CO}>$ K and densities greater than $_{\rm365 H} > 10^{10}$ $^{-3}$; but temperatures should also be less than K, otherwise CO molecules would dissociate \citep{scoville1980}."366 Then. we convolved the absorption cocficicut as a function of wavoleugth with the velocity profile of a disk in Weplerian rotation (see below).," Then, we convolved the absorption coefficient as a function of wavelength with the velocity profile of a disk in Keplerian rotation (see below)."367 We then multiplied the ‘absorption] cocfficieut bv the column density of CO in order to obtain the optical depth., We then multiplied the absorption coefficient by the column density of CO in order to obtain the optical depth.368 From the optical depth. we calculated line intensities using the transfer (∖⋠⋅()u ⋅ oe Oi?. .," From the optical depth, we calculated line intensities using the transfer equation (Eqn."369" . quan»COCLOC ""iQ hoe!""""OC|""CO", 9 in \citealt{kraus2000}) ).370" isod iTOIOClocitht ‘heOLLIC&le OTof a ""CLISISdiskps atj KelSCDACHAMLCDICYE FOTATION""Olet() TH TOCOderOYOCI TO inCI ο1ΝοTVATIOMΝεOUS O the CO baudhead emission.", \citet{aspin2010} needed to introduce a velocity profile of a disk in Keplerian rotation in order to fit their observations of the CO bandhead emission.371" They found that the best fit could be obtained using a stellar mass of MAL... inner disk radius of 13,=0.08 AAU. outer disk radius of AAT. temperature of Teo=2500 and optical r,4,20.13depth του<0.1."," They found that the best fit could be obtained using a stellar mass of $_*$ $_{\odot}$, inner disk radius of $_{\rm in}$ AU, outer disk radius of $_{\rm out}$ AU, temperature of $_{\rm CO}$ K and optical depth $\tau_{\rm372 CO}<0.1$."373 We calculated the line KINprofile clucreine from such a disk. assuuiug a disk inclination of and a disk radial brightuess profile proportional to 2? (same as in ?)).," We calculated the line profile emerging from such a disk, assuming a disk inclination of $^{\circ}$ and a disk radial brightness profile proportional to $^{-2.5}$ (same as in \citealt{goto2011}) )."374 The resulting line profile is otted in the small inset in Fie. 7.., The resulting line profile is plotted in the small inset in Fig. \ref{fig:co}.375 We used this profile when calculating the optical depth (sce above)., We used this profile when calculating the optical depth (see above).376 Finally. we smoothed our model spectru to the instrmucutal resolution of SINFONT.," Finally, we smoothed our model spectrum to the instrumental resolution of SINFONI."377 The result. plotted with a grav ino in Fie. 7..," The result, plotted with a gray line in Fig. \ref{fig:co},"378 fits our observations of the 20. »l.aud v=L >2 transitions very well.," fits our observations of the $\rightarrow$ 0, $\rightarrow$ 1, and $\rightarrow$ 2 transitions very well."379" We note that a 1n0doel with au outer radius of r4420.1 AAU, as derived wOT from the profiles of the CO fundamental cimission ines of LLup. would ft our observations equally well."," We note that a model with an outer radius of $_{\rm out}$ AU, as derived by \citet{goto2011} from the profiles of the CO fundamental emission lines of Lup, would fit our observations equally well."380 ⋅ - Ca. Ik. Fe. Ti. aud. Si absorption lines iu their near-infrared spectra. which are often used for spectral type classificationf: E(2)..," M-type stars display numerous Na, Ca, K, Fe, Ti, and Si absorption lines in their near-infrared spectra, which are often used for spectral type classification \citep{cushing2005}."381 Tn quiescence.: LEup also displayed. atoutic: absorption: lines: (ST at 1.199;12. MgILI at μια. and ALIE at 1.67240. sce E 23).," In quiescence, Lup also displayed atomic absorption lines I at $\,\mu$ m, I at $\,\mu$ m, and I at $\,\mu$ m, see \citealt{sipos2009}) )."382 These1 lines. are most probably photospheric iu origin (seealso?).., These lines are most probably photospheric in origin \citep[see also][]{herbig2001}.383 During outburst. no absorption lines are visible any more. despite the fact that. at the time of our SINFONT observation. of the J aud T-band fux. and of the K-band fux was photospheric.," During outburst, no absorption lines are visible any more, despite the fact that, at the time of our SINFONI observation, of the J and H-band flux, and of the K-band flux was photospheric."384 These values were obtained by supposing that the quiescent| JITIEK photometry represent the stellar plotosphere. aud that duving our observations. the V...J. V. ιν and VIN colors," These values were obtained by supposing that the quiescent JHK photometry represent the stellar photosphere, and that during our observations, the $-$ J, $-$ H, and $-$ K colors"385and Z presented in Figs.,and $Z$ presented in Figs.386 8 and 9 is also found in these observations., \ref{fig:L-Gals_SFR-Z_SSFR-Z_z=0} and \ref{fig:sSFR-GasFrac} is also found in these observations.387" In this section, we study the origin of the turnover in the model M.- Z relation seen in Fig. 7.."," In this section, we study the origin of the turnover in the model $M_{*}$ $Z$ relation seen in Fig. \ref{fig:L-Gals_ave_z=0}."388 We do this by splitting the high-mass end of the sample into low metallicity and high metallicity sub-populations and studying differences in their evolutionary histories., We do this by splitting the high-mass end of the sample into low metallicity and high metallicity sub-populations and studying differences in their evolutionary histories.389" We extracted two high-mass (M,> 10!°°M.,) sub-samples: a high-Z (Z2 9.2) sub-sample containing 134 galaxies, and a low-Z (Z< 9.0) sub-sample containing 136 galaxies."," We extracted two high-mass $M_{*} > 10^{10.8}\textnormal{M}_{\textnormal{\astrosun}}$ ) sub-samples: a $Z$ $Z\geq9.2$ ) sub-sample containing 134 galaxies, and a $Z$ $Z\leq9.0$ ) sub-sample containing 136 galaxies."390" The mass, metallicity and SFR evolution of these two sub-samples was then compared."," The mass, metallicity and SFR evolution of these two sub-samples was then compared."391" The model tracks six distinct mass components of galaxies: stellar mass (in the form of a bulge and disc), black hole mass, cold gas mass (ISM), hot gas mass (ICM), ejected gas mass (IGM) and halo stars (producing the intra-cluster light)."," The model tracks six distinct mass components of galaxies: stellar mass (in the form of a bulge and disc), black hole mass, cold gas mass (ISM), hot gas mass (ICM), ejected gas mass (IGM) and halo stars (producing the intra-cluster light)."392" Mass and metals can pass between these components along pre-defined routes, depending on the processes taking place."," Mass and metals can pass between these components along pre-defined routes, depending on the processes taking place."393 The top three panels of Fig., The top three panels of Fig.394 10 show the time evolution of these mass components for three representative galaxies from the model with high stellar masses and high metallicities., \ref{fig:highZ_specgals} show the time evolution of these mass components for three representative galaxies from the model with high stellar masses and high metallicities.395" In the middle panels, we show the time evolution of the stellar and gas-phase metallicities of the same galaxies."," In the middle panels, we show the time evolution of the stellar and gas-phase metallicities of the same galaxies."396 The bottom panels show the time evolution of their star formation rates., The bottom panels show the time evolution of their star formation rates.397 The left panels display a type 0 galaxy in which the stellar mass (solid orange line) has been steadily increasing since redshift two., The left panels display a type 0 galaxy in which the stellar mass (solid orange line) has been steadily increasing since redshift two.398" In this galaxy, the mass of cold gas (solid blue line) is always higher than the critical value required for star formation, Mat (dashed blue There is a steady, gradual increase in the metallicity of the stellar, cold gas and hot gas components."," In this galaxy, the mass of cold gas (solid blue line) is always higher than the critical value required for star formation, $M_{\textnormal{crit}}$ (dashed blue There is a steady, gradual increase in the metallicity of the stellar, cold gas and hot gas components."399" This is because stars are formed continuously, synthesising and distributing metals throughout the galaxy at a higher rate than the dilution due to the accretion of metal-poor gas."," This is because stars are formed continuously, synthesising and distributing metals throughout the galaxy at a higher rate than the dilution due to the accretion of metal-poor gas."400 Around 64 per cent of the galaxies in our high-Z sub-sample have formation histories similar to this., Around 64 per cent of the galaxies in our $Z$ sub-sample have formation histories similar to this.401seeing is estimated to be twice more frequent at 20m. 20 times versus 10 times per winter. and that is also easy to understand with à reasonable assumption on the vertical motions of the boundary layer upper limit.,"seeing is estimated to be twice more frequent at 20m, 20 times versus 10 times per winter, and that is also easy to understand with a reasonable assumption on the vertical motions of the boundary layer upper limit."402of DM particles in the vicinity of the Earth is relevant for direct detection experiments.,of DM particles in the vicinity of the Earth is relevant for direct detection experiments.403" Vogelsbergeretal.(2009) showed that although this velocity distribution is quite well approximated by a smooth trivariate Gaussian, potentially measurable features are imprinted on the corresponding energy distribution by the detailed formation history of the Milky Way's halo."," \cite{Vogelsberger2009} showed that although this velocity distribution is quite well approximated by a smooth trivariate Gaussian, potentially measurable features are imprinted on the corresponding energy distribution by the detailed formation history of the Milky Way's halo."404 Here we concentrate on the velocities of the different phase-space components in the inner halo., Here we concentrate on the velocities of the different phase-space components in the inner halo.405 In Fig., In Fig.406" 10 we show v,-v projections of the distribution of all particles in the radial range 6 to 12 kpc (top), of those in substructures (middle), and of those in self-bound subhaloes (bottom)."," \ref{vrvt} we show $v_r$ $v_t$ projections of the distribution of all particles in the radial range 6 to 12 kpc (top), of those in substructures (middle), and of those in self-bound subhaloes (bottom)."407" These plots are two-dimensional histograms, with colour encoding the mass in the corresponding bin as indicated by the colour bar (in units of solar masses per 2 km/s x 2 km/s pixel)."," These plots are two-dimensional histograms, with colour encoding the mass in the corresponding bin as indicated by the colour bar (in units of solar masses per $2$ km/s x $2$ km/s pixel)."408 The total mass contributing to each panel is given in its top right-hand corner., The total mass contributing to each panel is given in its top right-hand corner.409" The main halo and so the bulk of the particles lie primarily at velocities below 200 km s!, whereas subhaloes and tidal streams are found almost exclusively at higher velocities."," The main halo and so the bulk of the particles lie primarily at velocities below 200 km $^{-1}$, whereas subhaloes and tidal streams are found almost exclusively at higher velocities."410" As a result, the most massive subhaloes are still (just) visible in the top panel despite the fact that they contribute less than a tenth of a percent of the mass."," As a result, the most massive subhaloes are still (just) visible in the top panel despite the fact that they contribute less than a tenth of a percent of the mass."411" These structures contribute to the high-energy tail of the recoil spectrum in direct DM detection experiments, and so may be visible in high resolution experiments, particularly those with directional sensitivity which can detect the common motion of the substructure particles."," These structures contribute to the high-energy tail of the recoil spectrum in direct DM detection experiments, and so may be visible in high resolution experiments, particularly those with directional sensitivity which can detect the common motion of the substructure particles."412 We study the population of subhaloes and tidal streams in six Milky Way-like DM haloes taken from the Aquarius Project., We study the population of subhaloes and tidal streams in six Milky Way-like DM haloes taken from the Aquarius Project.413" These structures are identified using the Hierarchical Structure Finder (HSFMaciejewskietal. 2009a),, a state-of-the-art structure finder which operates in 6-D phase-space."," These structures are identified using the Hierarchical Structure Finder \citep[HSF][]{Maciejewski2009a}, a state-of-the-art structure finder which operates in 6-D phase-space."414 We find that that the differential mass function of self-bound subhaloes can be well described by a power-law with slope close to —1.9., We find that that the differential mass function of self-bound subhaloes can be well described by a power-law with slope close to $-1.9$.415 This agrees with results from an independent analysis using the 3-D structure finder SUBFIND (Springeletal.2008)., This agrees with results from an independent analysis using the 3-D structure finder SUBFIND \citep{Springel2008}.416". Typically HSF attaches slightly more particles to subhaloes than SUBFIND, and also finds slightly more subhaloes above the simulation resolution limit (see Table 1))."," Typically HSF attaches slightly more particles to subhaloes than SUBFIND, and also finds slightly more subhaloes above the simulation resolution limit (see Table \ref{table2}) )."417 This agrees with previous results described in Maciejewskietal.(2009a)., This agrees with previous results described in \cite{Maciejewski2009a}.418". About of the mass within rso is in self-bound subhaloes, with significant scatter among the six haloes Aq-A to Aq-F. HSF subhalo masses are ~10% larger than those found by SUBFIND, although the increase can be larger near halo centre."," About of the mass within $r_{50}$ is in self-bound subhaloes, with significant scatter among the six haloes Aq-A to Aq-F. HSF subhalo masses are $\sim 10\%$ larger than those found by SUBFIND, although the increase can be larger near halo centre."419 In most haloes the total subhalo mass is dominated by the largest objects., In most haloes the total subhalo mass is dominated by the largest objects.420" The radial distributions of HSF and SUBFIND subhaloes are almost identical, although HSF can identify subhaloes closer to halo centre due to their enhanced density contrast in phase-space."," The radial distributions of HSF and SUBFIND subhaloes are almost identical, although HSF can identify subhaloes closer to halo centre due to their enhanced density contrast in phase-space."421" The differential mass function for substructures (i.e. both subhaloes and tidal streams) is also well described with a power-law, but in this case the slope is close to —2."," The differential mass function for substructures (i.e. both subhaloes and tidal streams) is also well described with a power-law, but in this case the slope is close to $-2$."422" This is independent of simulation resolution and holds approximately for all six haloes, thus appearing robust."," This is independent of simulation resolution and holds approximately for all six haloes, thus appearing robust."423 For most of the level 2 haloes around 35% of the mass within rso is assigned to substructures with mass above, For most of the level 2 haloes around $35\%$ of the mass within $r_{50}$ is assigned to substructures with mass above424 The simulations used in (his study were performed withEnzo. au Everian adaptive mesh-relinement (AMIR). byerodynamical + N-body eode (Bryan & Norman 1997: O'Sheaelal. 2004. 2005).," The simulations used in this study were performed with, an Eulerian adaptive mesh-refinement (AMR), hydrodynamical + N-body code (Bryan & Norman 1997; O'Shea 2004, 2005)."425 Smithetal. (2011) enhanced this code by adding new modules for star formation. primordial chemistry. and cooling rates consistent wilh ionizing radiation. metal transport. and feedback.," Smith (2011) enhanced this code by adding new modules for star formation, primordial chemistry, and cooling rates consistent with ionizing radiation, metal transport, and feedback."426 The ionizing background is spatially constant. ancl optically. thin. but variable in redshift.," The ionizing background is spatially constant and optically thin, but variable in redshift."427 At this stage. we have not implemented radiative transfer or spectral filtering bv the IGM.," At this stage, we have not implemented radiative transfer or spectral filtering by the IGM."428 For the chunping factor caleulation. we ran a simulation on a 50/7.| Mpe static grid (unierid) cube with 1024* cells. denoted as run 5010242 in Table 1 of Smithetal. (2011).," For the clumping factor calculation, we ran a simulation on a $h^{-1}$ Mpc static grid (“unigrid"") cube with $1024^3$ cells, denoted as run 50_1024_2 in Table 1 of Smith (2011)."429 The radiative heating from the ionizing background plavs an important role in determining the properties of the filamentary structure., The radiative heating from the ionizing background plays an important role in determining the properties of the filamentary structure.430 As a filament is ionizecl and heated. its density drops and its temperature rises: both effects reduce the recombination rate.," As a filament is ionized and heated, its density drops and its temperature rises; both effects reduce the recombination rate."431 To study these effects. we ran four moderate-resolution (505.+ Mpe unigrid cube with 512* cells) simulations with varving ionizing backgrounds. summarized in Table 2.," To study these effects, we ran four moderate-resolution $h^{-1}$ Mpc unigrid cube with $512^3$ cells) simulations with varying ionizing backgrounds, summarized in Table 2."432" After initial submission of this paper. we ran a 15367 simulation. to check convergence and assess (he “cosmic variance"" among eight sub-volumes of the 1024* and 1536* simulations."," After initial submission of this paper, we ran a $1536^3$ simulation, to check convergence and assess the “cosmic variance"" among eight sub-volumes of the $1024^3$ and $1536^3$ simulations."433 The standard UV background was taken from Iaardt & Aladau (2001). allhough we also explore new computations of high-z SFRs by Trentietal. (2010) ancl Haardt & Macau (2012).," The standard UV background was taken from Haardt & Madau (2001), although we also explore new computations of $z$ SFRs by Trenti (2010) and Haardt & Madau (2012)."434 These four simulations were: (1) no photoionizing background: (2) UV background ramped up from 2=7 lo 2=G6 (run 505122 from Smithetal. 2011): (2) UV background ramped up from 2:=9 to 2—S and (4) UV background ramped up from z=9 to z=& but twice as strong as in (3).," These four simulations were: (1) no photoionizing background; (2) UV background ramped up from $z=7$ to $z=6$ (run 50_512_2 from Smith 2011); (3) UV background ramped up from $z=9$ to $z=8$; and (4) UV background ramped up from $z=9$ to $z=8$, but twice as strong as in (3)."435 Post-processing of the simulations was performed using the data analvsis aud visualization package.yt!. documented by Turketal. For regions of ionized hvdrogen of density nyyy. we caleulate the chuupine factor. Cy.," Post-processing of the simulations was performed using the data analysis and visualization package, documented by Turk For regions of ionized hydrogen of density $n_{\rm HII}$, we calculate the clumping factor, $C_H$."436 using (wo clifferent methods., using two different methods.437 The first. caleulation. which has been used in some earlier studies. uses density. weighting.," The first calculation, which has been used in some earlier studies, uses density weighting."438 In this “density field” (DE) method we define," In this “density field"" (DF) method we define"439As can be seen from Tabs.,As can be seen from Tabs.440 3. and 4.. the low photon statistics do not allow the temperature to be constrained. although 1f we assume an equipartition between the galaxy velocities and the gas. Le. AT=HIIoe. we derive a temperature of AT=3.9+1.5 keV —- in excellent agreement with the best fit value.," \ref{tab:x1} and \ref{tab:x2}, the low photon statistics do not allow the temperature to be constrained, although if we assume an equipartition between the galaxy velocities and the gas, i.e. $kT = \mu m_p \sigma^{2}_v$, we derive a temperature of $kT=3.9\pm1.5$ keV – in excellent agreement with the best fit value."441 Morphologically ((Fig. 4)), Morphologically (Fig. \ref{fig:clsa:image}) )442 looks like present day still relaxing clusters: there is an obvious centre formed by the brightest cluster galaxies displaced at ~10” from the X-ray emission peak., looks like present day still relaxing clusters: there is an obvious centre formed by the brightest cluster galaxies displaced at $\sim 10\arcsec$ from the X-ray emission peak.443 The brightest cluster galaxy (BCG) has a typical giant elliptical galaxy absorption spectrum and is at 750 km s! with respect to the mean cluster redshift., The brightest cluster galaxy (BCG) has a typical giant elliptical galaxy absorption spectrum and is at 750 km $^{-1}$ with respect to the mean cluster redshift.444 The velocity distribution is quite broad. but the null hypothesis that the observed distribution is drawn from normal cannot be ruled out — the Anderson-Darling statistics value is 0.527 which is less than the critical value 0.616 needed to reject the normality assumption at confidence level.," The velocity distribution is quite broad, but the null hypothesis that the observed distribution is drawn from normal cannot be ruled out – the Anderson-Darling statistics value is 0.527 which is less than the critical value 0.616 needed to reject the normality assumption at confidence level."445 There are two other X-ray sources projected over the cluster emission., There are two other X-ray sources projected over the cluster emission.446 One is the subclump to the north-east. which is most likely associated with the cluster as there are three cluster members within the X-ray contours (two of which are emission line galaxies) and no obvious optical counterpart at the X-ray peak.," One is the subclump to the north-east, which is most likely associated with the cluster as there are three cluster members within the X-ray contours (two of which are emission line galaxies) and no obvious optical counterpart at the X-ray peak."447 Unfortunately. due to instrumental problems. we could not get the optical spectrum of the south-east point-like source.," Unfortunately, due to instrumental problems, we could not get the optical spectrum of the south-east point-like source."448 The X-ray characteristics of the cluster are relatively well constrained as we have about 1600 photons in the spectrum and the cluster is of low temperature., The X-ray characteristics of the cluster are relatively well constrained as we have about 1600 photons in the spectrum and the cluster is of low temperature.449 The low luminosity of the cluster is 1n agreement with its temperature and the velocity dispersion. as compared to the scaling relations of local clusters (Fig. 3)). (," The low luminosity of the cluster is in agreement with its temperature and the velocity dispersion, as compared to the scaling relations of local clusters (Fig. \ref{fig:scaling}) ). ("450(Fig. 5)),Fig. \ref{fig:clsb:image}) )451" resembles apoint-like source in X-rays. although its half-light radius is Rsyj=17” (see Table 3.. to be compared to the PSF half-energy radius of 9"") and the extension log-likelihood (as reported by the XMM-SAS task emldetect) is 7.4. corresponding to probability of 6x107 of being a point-like source."," resembles apoint-like source in X-rays, although its half-light radius is $R_{50} \approx17\arcsec$ (see Table \ref{tab:x1}, to be compared to the PSF half-energy radius of $9\arcsec$ ) and the extension log-likelihood (as reported by the XMM-SAS task ) is 7.4, corresponding to probability of $6\times10^{-4}$ of being a point-like source."452 The 8 centrally located galaxies form a crown-like figure. and five are in a very narrow redshift range 0.77120.775 most likely forming a compact core.," The 8 centrally located galaxies form a crown-like figure, and five are in a very narrow redshift range $0.771-0.775$ most likely forming a compact core."453 The brightest cluster galaxy is at redshift coinciding with the mean of the centrally located galaxies (zpec;;= 0.7714) and has a typical giant elliptical galaxy absorption spectrum., The brightest cluster galaxy is at redshift coinciding with the mean of the centrally located galaxies $z_{BCG}=0.7714$ ) and has a typical giant elliptical galaxy absorption spectrum.454 We have to note that there are galaxies at similar redshifts (z= 0.77) all over the field as can be seen in Fig. 5..," We have to note that there are galaxies at similar redshifts $z455\approx 0.77$ ) all over the field as can be seen in Fig. \ref{fig:clsb:image}."456 None of the other X-ray sources present in the field are cluster members., None of the other X-ray sources present in the field are cluster members.457 The temperature is quite tightly constrained by the X- spectral fit even though the photon statistics seems rather poor — 480 photons in the [0.2-10] keV band in total., The temperature is quite tightly constrained by the X-ray spectral fit even though the photon statistics seems rather poor – 480 photons in the [0.2-10] keV band in total.458 This is possible as the cluster is at low temperature., This is possible as the cluster is at low temperature.459 The relations between cluster global characteristics. Ly.T and oy (Fig. 3))," The relations between cluster global characteristics, $L_X,\ T$ and $\sigma_V$ (Fig. \ref{fig:scaling}) )"460 are consistent with those observed from local clusters., are consistent with those observed from local clusters.461 There ts a large fraction of emission-line galaxies in the centre (5 from 8) as identified by the presence of the [OI] 3727 lline., There is a large fraction of emission-line galaxies in the centre (5 from 8) as identified by the presence of the [OII] 3727 line.462" This line is not broad in Type | AGN. and thus it is difficult to assess the presence of an AGN with only this one emission feature — most of the diagnostic line ratios that can be used involve lines (77. [OIII] 4959+5007A. H,) outside the range of the grism RIG00 used in this spectroscopic run."," This line is not broad in Type 1 AGN, and thus it is difficult to assess the presence of an AGN with only this one emission feature – most of the diagnostic line ratios that can be used involve lines $H_{\beta}$, [OIII] 4959+5007, $H_{\alpha}$ ) outside the range of the grism RI600 used in this spectroscopic run."463 Optical-to-X-ray flux ratios from deep surveys suggest that AGN with magnitudes typical of these galaxies have X-ray fluxes covering a wide range. between z8x 107? and «8x107  erg em s! (0.5-2 keV). with the mean being =2x 107) erg em7 Ss”! (eg.," Optical-to-X-ray flux ratios from deep surveys suggest that AGN with magnitudes typical of these galaxies have X-ray fluxes covering a wide range, between $\approx$ $^{-16}$ and $\approx$ $^{-15}$ erg $^{-2}$ $^{-1}$ (0.5-2 keV), with the mean being $\approx$ $10^{-15}$ erg $^{-2}$ $^{-1}$ (eg."464 Mainieri et al. 2002))., Mainieri et al. \cite{mai02}) ).465 Even in the unlikely event that all 5 emission line galaxies contain X-ray emitting AGN at this mean flux. the total would only be «θές of the measured cluster flux.," Even in the unlikely event that all 5 emission line galaxies contain X-ray emitting AGN at this mean flux, the total would only be $\approx$ of the measured cluster flux."466 It is perhaps more likely that most are Butcher-Oemler star-forming galaxies. implying an interestingly high star-forming fraction (eg.," It is perhaps more likely that most are Butcher-Oemler star-forming galaxies, implying an interestingly high star-forming fraction (eg."467 Fairley et al., Fairley et al.468 2002 find blue fractions up to in low Ly clusters at z~0.3)., \cite{fai02} find blue fractions up to in low $L_X$ clusters at $\sim$ 0.3).469 Pending further studies of the galaxy population. we conclude that some contamination by AGN may be possible. but that it ts unlikely that all the flux originates in AGN.," Pending further studies of the galaxy population, we conclude that some contamination by AGN may be possible, but that it is unlikely that all the X-ray flux originates in AGN."470 Based on the morphology and the redshift distribution. the most plausible interpretation is that the cluster is in its early stage of formation: a compact core is already formed and the aceretion of matter from the nearby large-scale structure Is underway — many galaxies in the field have redshifts around the redshift of the central core.," Based on the morphology and the redshift distribution, the most plausible interpretation is that the cluster is in its early stage of formation: a compact core is already formed and the accretion of matter from the nearby large-scale structure is underway – many galaxies in the field have redshifts around the redshift of the central core."471 The cluster X-ray emission (see Fig. 7)), The cluster X-ray emission (see Fig. \ref{fig:clsd:image}) )472 lies between two bright X-ray sources identified as QSOs at z=1.12 (the closest to the north-west. with 486 counts in the [0.5-5|] keV band and V = 21.2. R = 20.9. I = 20.8) and at zc1.19 (to the south-east. with 378 counts and V = 18.5. R = 18.5. 1 = 18.0).," lies between two bright X-ray sources identified as QSOs at $z = 1.12$ (the closest to the north-west, with 486 counts in the [0.5-5] keV band and V = 21.2, R = 20.9, I = 20.8) and at $z = 1.19$ (to the south-east, with 378 counts and V = 18.5, R = 18.5, I = 18.0)."473 These two QSOs are located in diametrically opposed directions at 0.9 and 1.2’ from the centre of the cluster X-ray emission. respectively.," These two QSOs are located in diametrically opposed directions at $0.9\arcmin$ and $1.2\arcmin$ from the centre of the cluster X-ray emission, respectively."474 Such à configuration 1s very unlikely to occur by chance., Such a configuration is very unlikely to occur by chance.475 Our estimated surface density is approximately 30 X-ray sources per deg over a field of 2.8 deg? in the XMM-LSS survey with detected counts greater than 300. subsequently the probability of detecting two X-ray sources within 2’ from the cluster X-ray centre is about1%.," Our estimated surface density is approximately 30 X-ray sources per $^2$ over a field of 2.8 $^2$ in the XMM-LSS survey with detected counts greater than 300, subsequently the probability of detecting two X-ray sources within $2\arcmin$ from the cluster X-ray centre is about."476. Given that these two quasars are located just behind a putative foreground cluster at z = 0.9 (see below). gravitational lensing could possibly affect the real positions and fluxes of these two background objects.," Given that these two quasars are located just behind a putative foreground cluster at z = 0.9 (see below), gravitational lensing could possibly affect the real positions and fluxes of these two background objects."477 However. the hypothesis that these QSOs could be two images of a single source lensed by a foreground cluster does not hold as their optical spectra show a significant redshift difference Az5 0.07.," However, the hypothesis that these QSOs could be two images of a single source lensed by a foreground cluster does not hold as their optical spectra show a significant redshift difference $\Delta z478\simeq 0.07$ ."479 A more detailed discussion on this unusual association will be presented elsewhere (Jean et al..," A more detailed discussion on this unusual association will be presented elsewhere (Jean et al.,"480 in preparation)., in preparation).481iowever. the relative suppression of the star-forming halos increases faster with redshift. and eventually it becomes arger than that of the minihalos (beyond. 2~50).,"however, the relative suppression of the star-forming halos increases faster with redshift, and eventually it becomes larger than that of the minihalos (beyond $z482\sim 50$ )."483 At ;=20. the bulk velocities reduce the mean eas fraction in star-forming halos by a [actor of 1.8 and that in minihalos w 3.1.," At $z=20$, the bulk velocities reduce the mean gas fraction in star-forming halos by a factor of 1.8 and that in minihalos by 3.1."484 Unlike previous studies. in our calculations the relative velocities produce three separate effects: suppression of the ido abundance (dnidÀL). suppression of the gas content within cach halo (ΑΛ) ). ancl boosting of the minimum cooling mass (through μη) ).," Unlike previous studies, in our calculations the relative velocities produce three separate effects: suppression of the halo abundance $dn/dM$ ), suppression of the gas content within each halo $f_g(M)\,$ ), and boosting of the minimum cooling mass (through $V\cool(z)\,$ )."485 In order to gain a better physical understanding. and for easier comparison. with oevious papers. we investigate the relative importance of each elfect in Figure 3..," In order to gain a better physical understanding, and for easier comparison with previous papers, we investigate the relative importance of each effect in Figure \ref{Fig:Vc2rat}."486 For the star-forming halos. the suppression of gas content is always the least. significant cect (e.g. suppression by a factor of 1.13 on its own at >= 20). while the cooling mass boost is most. important above z=28.5 (factor of 1.26 on its own at z= 20). and the halo abundance cut is most important at lower redshifts (factor of 1.43 on its own at 2= 20).," For the star-forming halos, the suppression of gas content is always the least significant effect (e.g., suppression by a factor of 1.13 on its own at $z=20$ ), while the cooling mass boost is most important above $z=28.5$ (factor of 1.26 on its own at $z=20$ ), and the halo abundance cut is most important at lower redshifts (factor of 1.43 on its own at $z=20$ )."487 For the minihalos. the boosting of the minimum cooling mass acts as a (small) positive effect. since it moves eas from the star-forming to the minihalo category (0... boost by a factor of 1.10 on its own at 2= 20). while the other two elfects are larger ancl comparable (e.g.. at z=20 the suppression of gas content would. give a reduction bx a factor of 2.17 on its own. and the halo abundance cut would. give a suppression factor of 1.14).," For the minihalos, the boosting of the minimum cooling mass acts as a (small) positive effect, since it moves gas from the star-forming to the minihalo category (e.g., boost by a factor of 1.10 on its own at $z=20$ ), while the other two effects are larger and comparable (e.g., at $z=20$ the suppression of gas content would give a reduction by a factor of 2.17 on its own, and the halo abundance cut would give a suppression factor of 1.74)."488 The gas fractions shown in Figures 2 and 3 are globally averaged., The gas fractions shown in Figures 2 and 3 are globally averaged.489 Llowever. in reality. the universe ds. highly inhomogencous on small cosmological scales.," However, in reality the universe is highly inhomogeneous on small cosmological scales."490 We can clivice it into patches that have various bulk velocities and densities., We can divide it into patches that have various bulk velocities and densities.491 In this section we consider just the variation with velocity. ie. averaged over all density Huctuations.," In this section we consider just the variation with velocity, i.e., averaged over all density fluctuations."492 In other words. we look at the contribution of velocity Ductuations to [luctuations in the gas fraction in halos.," In other words, we look at the contribution of velocity fluctuations to fluctuations in the gas fraction in halos."493 Lf we consider patches that are still small enough to have a coherent. (y (e.g. cubes of 3 comoving Alpe on a side). then the absolute value of the bulk velocity in each one follows a Maxwell-3oltzmann distribution (eq. 1)).," If we consider patches that are still small enough to have a coherent $v\bc$ (e.g., cubes of 3 comoving Mpc on a side), then the absolute value of the bulk velocity in each one follows a Maxwell-Boltzmann distribution (eq. \ref{MBdist}) )."494 Consider. the contributions of patehes of various velocities to the total amount of star formation., Consider the contributions of patches of various velocities to the total amount of star formation.495 At a given redshift. the gas fraction in star-forming halos is lower in the xuehes with a high value of the relative velocity. because all hree velocity cllects (see the previous subsection) tend to reduce this gas fraction.," At a given redshift, the gas fraction in star-forming halos is lower in the patches with a high value of the relative velocity, because all three velocity effects (see the previous subsection) tend to reduce this gas fraction."496 On the other hand. patches with zero bulk velocity do not contribute much. simply because hey are rare.," On the other hand, patches with zero bulk velocity do not contribute much, simply because they are rare."497" As shown in the top panel of Figure 4.. he most common bulk velocity is íi;~O0S20,,. where me and ao, are both measured at. the same. recshilt (recombination or anv other 2)."," As shown in the top panel of Figure \ref{Fig:VcPdf}, the most common bulk velocity is $v\bc \sim4980.82 \sigma_{v\bc}$, where $v\bc$ and $\sigma_{v\bc}$ are both measured at the same redshift (recombination or any other $z$ )."499 Wf the stellar density. were independent. of the bulk velocity. then the contribution of regions of various velocities would be proportional to the velocity PDE.," If the stellar density were independent of the bulk velocity, then the contribution of regions of various velocities would be proportional to the velocity PDF."500 Instead. the velocity suppression effect shifts the contribution to stellar density (assumed: proportional to the gas [fraction in star-forming halos) towards lower me. with the relative change (compared to the Maxwell)oltzmann distribution) increasing strongly with recshilt.," Instead, the velocity suppression effect shifts the contribution to stellar density (assumed proportional to the gas fraction in star-forming halos) towards lower $v\bc$ , with the relative change (compared to the Maxwell-Boltzmann distribution) increasing strongly with redshift."501" Thus. the biggest contribution to stellar density comes from Due=0.670, patches at z=20. and from (i;=0.2304. patches at >=60."," Thus, the biggest contribution to stellar density comes from $v\bc = 0.67 \sigma_{v\bc}$ patches at $z=20$, and from $v\bc = 0.23 \sigma_{v\bc}$ patches at $z=60$."502 We compare the contributions of the three separate collects of the velocity to the shift in the distribution of star formation (Figure 4.. bottom panel).," We compare the contributions of the three separate effects of the velocity to the shift in the distribution of star formation (Figure \ref{Fig:VcPdf}, , bottom panel)."503 Xs in the top panel of Figure 3.. we find that the suppression of halo gas content has the least significant ellect on star-forming halos ab 2=20 (tvpically. à ~10% elfect on the distribution). while the other two effects (halo abundance suppression and cooling mass boost) have à ~20.—30% οσοι cach.," As in the top panel of Figure \ref{Fig:Vc2rat}, we find that the suppression of halo gas content has the least significant effect on star-forming halos at $z=20$ (typically, a $\sim 10\%$ effect on the distribution), while the other two effects (halo abundance suppression and cooling mass boost) have a $\sim 20-30\%$ effect each."504 Thus. at the highest. redshifts. the star formation is concentrated in low-velocity regions which are rare. ie. at the low-probahility ef. end. of the MaxwellBoltzmann distribution function.," Thus, at the highest redshifts, the star formation is concentrated in low-velocity regions which are rare, i.e., at the low-probability $v\bc^2$ end of the Maxwell-Boltzmann distribution function."505 Phe universe at these epochs is very inhomogeneous. with a few bright regions filled. with stars. while in all other regions the relative velocity is too high to allow significant star formation.," The universe at these epochs is very inhomogeneous, with a few bright regions filled with stars, while in all other regions the relative velocity is too high to allow significant star formation."506 As the universe expands. the relative velocity decays. and in more and more patches across the universe the relative velocity drops enough. to allow for star formation.," As the universe expands, the relative velocity decays, and in more and more patches across the universe the relative velocity drops enough to allow for star formation."507 As a result. the stellar clistributionbecomes increasingly homogeneous.," As a result, the stellar distributionbecomes increasingly homogeneous."508 Toquantify the degree of inhomogeneity caused by the dependence of stellar density, Toquantify the degree of inhomogeneity caused by the dependence of stellar density509efforts have been made since then to detect such changes using photospheric magnetic field observations.,efforts have been made since then to detect such changes using photospheric magnetic field observations.510 Several (vpes of changes in (he measured magnetic fields [rom pre- ancl post- of flares have been reported by ground-based (Patterson&Zirin19381:PattersonWangetal.1992:Ambastha1993:Chen1994:Tagvard 1999).. and more recently by space-based (Ixosovichev&Zharkova2001:QiuGary2003:Wane2006:AMaurva&Ambastha2008.2009:Marva2010) observations.," Several types of changes in the measured magnetic fields from pre- and post-phases of flares have been reported by ground-based \citep{Patterson1981,511Patterson1984, Wang1992, Ambastha1993, Chen1994, Hagyard1999}, and more recently by space-based \citep{Kosovichev2001, Qiu2003,512Wang2006, Maurya2008, Maurya2009, Maurya2010} observations."513" These changes in observed magnetic fields can be divided into (wo main categories: “permanent” and ""transient.", These changes in observed magnetic fields can be divided into two main categories: “permanent” and “transient”.514 The first (ype is irreversible change observed. during pre- to post- phases of flares., The first type is irreversible change observed during pre- to post- phases of flares.515" The later. ""(ransient changes are reversible Chat occur only during the peak or impulsive phase οἱ energetic events."," The later, “transient” changes are reversible that occur only during the peak or impulsive phase of energetic events."516 The observed changes in magnetic field parameters during the impulsive phase of a flare are. however. expected to be affected by a variety of effects that could introduce ambiguity in the cause and effects of the observed changes.," The observed changes in magnetic field parameters during the impulsive phase of a flare are, however, expected to be affected by a variety of effects that could introduce ambiguity in the cause and effects of the observed changes."517 One major concern is that (the flare associated modification of spectral line profiles. used [or the measurement of photospheric magnetic fields. πια lead (o incorrect estimation of magnetic [fields during the inipulsive phase of major flares.," One major concern is that the flare associated modification of spectral line profiles, used for the measurement of photospheric magnetic fields, may lead to incorrect estimation of magnetic fields during the impulsive phase of major flares."518" Therelore. adequate care must be taken in interpreting (he ""observed"" changes in photospheric magnetic fields and Doppler velocities."," Therefore, adequate care must be taken in interpreting the “observed” changes in photospheric magnetic fields and Doppler velocities."519" For instance. the irreversible changes observed after [lares may occur due to flux emergence/cancelation process and may be considered as ""real changes."," For instance, the irreversible changes observed after flares may occur due to flux emergence/cancelation process and may be considered as “real” changes."520 However. there is ambiewly in interpreting (he t(ransients. reversible changes observed curing some large flares as it has been difficult to ascertain whether (hese changes are real or artifacts.," However, there is ambiguity in interpreting the transients, reversible changes observed during some large flares as it has been difficult to ascertain whether these changes are real or artifacts."521" These “magnetic transients” are therefore termed also as ""magnetic anomalies (Patterson1984:Ixosovichev&Zharkova2003:Marva&Ambastha 2009)."," These “magnetic transients” are therefore termed also as “magnetic anomalies” \citep{Patterson1984, Kosovichev2001, Qiu2003, Maurya2009}."522 some nunmerical experiments have been carried out to confirm magnetic field changes due to a transient change of the spectral line profile from absorption to emission (MachadoVernazzaetal.1981:Ding&Fang1989:Dine2002:Edelman 2004).," Some numerical experiments have been carried out to confirm magnetic field changes due to a transient change of the spectral line profile from absorption to emission \citep{Machado1980, Vernazza1981,523Ding1989, Ding2002, Edelman2004}."524. Elfect of line profile change on magnetic field estimation was also reported by Abramenko&(2004) using spectrographic observations of six different photospheric absorption lines at the flare and quiel locations.," Effect of line profile change on magnetic field estimation was also reported by \citet{Abramenko2004}525 using spectrographic observations of six different photospheric absorption lines at the flare and quiet locations."526 Thev showed that the core of spectral profiles at the flare location was shallower as compared to that lor the quiet Sun. and less steep in the wings.," They showed that the core of spectral profiles at the flare location was shallower as compared to that for the quiet Sun, and less steep in the wings."527" Thev attributed the enhanced core emission to thermal heating of photosphere by the Πάνο, and the less steep slope near wings to the inhomogeneity of the photospheric magnetic fields."," They attributed the enhanced core emission to thermal heating of photosphere by the flare, and the less steep slope near wings to the inhomogeneity of the photospheric magnetic fields."528 This resulted in an underestimation of the magnetic field measurements by 18-25% at the flare locations., This resulted in an underestimation of the magnetic field measurements by $\%$ at the flare locations.529 Qiu&Gary(2003) have reported (hat sien reversals of magnetic polarity may also be related to non-thermal processes., \citet{Qiu2003} have reported that sign reversals of magnetic polarity may also be related to non-thermal processes.530 They found anomalous polarity reversals during a lavee X5.6 flare of 2001 April 6 at the location of strong hard X-ray. (INI) emission forming near cooler umbral//penunmbral regionse of the sunspots., They found anomalous polarity reversals during a large X5.6 flare of 2001 April 6 at the location of strong hard X-ray (HXR) emission forming near cooler umbral/penumbral regions of the sunspots.531 These ΠΧ sources are produced by, These HXR sources are produced by532Darsukova et al. (,Barsukova et al. (5332002). IKato (2003) and Corauskü et al. (,"2002), Kato (2003) and Goranskii et al. ("5342001). aud. optical spectroscopy by Munan otf. al. (,"2004), and optical spectroscopy by Munari et al. ("5352002a.c0). Corauski et al. (,"2002a,c), Goranskii et al. ("5362002). IXolev. et al. (,"2002), Kolev et al. ("5372002). Osiwala et al. (,"2002), Osiwala et al. ("5382002). Kipper et al. (,"2002), Kipper et al. ("5392001).,2004).540 Modeling of the outburst was presented by Better Marom (2003). Soker Tyleuda (2003) and Boschi Miinari (2001). while modeling of the Lelt-eclo expausion was performed by Boud et al. (," Modeling of the outburst was presented by Retter Marom (2003), Soker Tylenda (2003) and Boschi Munari (2004), while modeling of the light-echo expansion was performed by Bond et al. ("5412003). Sugerman (2003) and Tylenda (2001).,"2003), Sugerman (2003) and Tylenda (2004)."542 Conuuon to uanuv of these papers were the issues of reddening. distance ancl nature of the progenitor of Vasa Mon.," Common to many of these papers were the issues of reddening, distance and nature of the progenitor of V838 Mon."543 However. noue of these papers focused specifically on these topics. aud nuited their discussion to the derivation or adoptiou of approximate values based ouessined cucrey distribution of the outbursting component.," However, none of these papers focused specifically on these topics, and limited their discussion to the derivation or adoption of approximate values based on energy distribution of the outbursting component."544 The aia of the preseut paper is to attenipt an accurate derivation of reddening aud distance of Va3e Mon by careful comparison of different. tecliuiques and approaches that do not involve the outbursting component. as well as to derive by comparison with theoretical models the nature aud evolutionary state of V8S2J8 Mon prior to the onset of the outburst.," The aim of the present paper is to attempt an accurate derivation of reddening and distance of V838 Mon by careful comparison of different techniques and approaches that do not involve the outbursting component, as well as to derive by comparison with theoretical models the nature and evolutionary state of V838 Mon prior to the onset of the outburst."545 Optical and IR photometry used in this paper was obtained with CC'D-equipped USNO Flagstaff Station (NOFS) telescopes., Optical and IR photometry used in this paper was obtained with CCD-equipped USNO Flagstaff Station (NOFS) telescopes.546 The data were obtained with the 1.hu telescope. the JIIR/L/ with the 1.55m telescope.," The data were obtained with the 1.0m telescope, the $^\prime$ $^\prime$ with the 1.55m telescope."547 The optical photometry is strictly tied to the Landolt (1983. 1992) svstem of equatorial standards. while the iufrared photometry uses standard JIN” filters and differential measures with respec to local 2\LASS stars.," The optical photometry is strictly tied to the Landolt (1983, 1992) system of equatorial standards, while the infrared photometry uses standard JHK' filters and differential measures with respect to local 2MASS stars."548 The low resolution optical spectropliotomoetrv of Va3s Mon was obtained with the AFOSC|CCD oenager|spectrograph of the 1.5214. telescope operated im Asiago bv the Astronomical Observatory of Padova., The low resolution optical spectrophotometry of V838 Mon was obtained with the AFOSC+CCD imager+spectrograph of the 1.82m telescope operated in Asiago by the Astronomical Observatory of Padova.549 At the same telescope we secured high resolution spectra with the Echelle|CCD spectrograph., At the same telescope we secured high resolution spectra with the Echelle+CCD spectrograph.550 Other high resolution spectra of Va3e Mon were obtained with the Couce 4.vectrograpl|CCD of the Bozheu 21m telescope. with lanURAFFE at the 1.9311 SAAO telescope. with SARC at je οι TNC telescope. and with FEROS at the 2.2m ESO telescope in La Silla.," Other high resolution spectra of V838 Mon were obtained with the Coude spectrograph+CCD of the Rozhen 2m telescope, with GIRAFFE at the 1.93m SAAO telescope, with SARG at the 3.5m TNG telescope, and with FEROS at the 2.2m ESO telescope in La Silla."551 Further details will be provided low where necessary., Further details will be provided below where necessary.552 Figure 1 presents a review of the optical aud IR xhotonietrie evolution of Vass Mou. updated to the ead of the third season of visibility in May. 2001.," Figure 1 presents a review of the optical and IR photometric evolution of V838 Mon, updated to the end of the third season of visibility in May 2004."553 With he exception of some carly-time datapoiuts. the lieht-curves are entirely based on data from the USNO clescopes aud therefore they are highly consistent aud free roni systematic effects of color transformation between different local plotometric svstenis that would badly affect observations [fron different observatorics of au object with such extreme colors as V838 Mon.," With the exception of some early-time datapoints, the light-curves are entirely based on data from the USNO telescopes and therefore they are highly consistent and free from systematic effects of color transformation between different local photometric systems that would badly affect observations from different observatories of an object with such extreme colors as V838 Mon."554 Va3a Mon les close to galactic equator (b= 11°) and in the auti-center quadrant (22187)., V838 Mon lies close to galactic equator $b$ $^\circ$ ) and in the anti-center quadrant $l$ $^\circ$ ).555 I spite the line of sight crosses the Oriou. Perseus and Outer spiral arlus. inspection of the Palomar plates. of the 2\LASS. IRAS. URST-IIo maps and the Neckel and Nlare (1980) extinction charts do support an appareutlv smooth star counts distribution for several degrees around V838 Mon position.," I spite the line of sight crosses the Orion, Perseus and Outer spiral arms, inspection of the Palomar plates, of the 2MASS, IRAS, $\alpha$ maps and the Neckel and Klare (1980) extinction charts do support an apparently smooth star counts distribution for several degrees around V838 Mon position."556 Therefore in. principle. open clusters and field stars with accurate photometry and spectroscopic classification can be used to search for a relatiouship between reddeuig and distance i the direction of Voa3ea Mon.," Therefore in principle, open clusters and field stars with accurate photometry and spectroscopic classification can be used to search for a relationship between reddening and distance in the direction of V838 Mon."557 Within 5° from Vass Mou there are 13 open clusters in the catalog of, Within $^\circ$ from V838 Mon there are 13 open clusters in the catalog of558proto-planets. planetesimals. comets. and asteroids Irequently bombarced the Solar surface and were accreted.,"proto-planets, planetesimals, comets, and asteroids frequently bombarded the Solar surface and were accreted."559 This accretecl matter is deficient in hydrogen ancl helium. introducing nelal-rich material into the upper lavers of the Sun.," This accreted matter is deficient in hydrogen and helium, introducing metal-rich material into the upper layers of the Sun."560 Exactly how much metal-rich matter was accreted in (hese early stages anc for how long. has been speculated by several authors (Christensen-Dalseaard. Gough. Morgan L979. Jelferv et al.," Exactly how much metal-rich matter was accreted in these early stages and for how long, has been speculated by several authors (Christensen-Dalsgaard, Gough, Morgan 1979, Jeffery et al."561 1997). but. unfortunately. is ιοί well known.," 1997), but unfortunately is not well known."562" More recently, Murrav et al. ("," More recently, Murray et al. ("5632001) have searched a sample of 640 solu-t(vpe stars or the signature of iron enhancement in (heir spectra and concluded that. on average. these stars appear (o have accreted about of iron while on the main sequence.,"2001) have searched a sample of 640 solar-type stars for the signature of iron enhancement in their spectra and concluded that, on average, these stars appear to have accreted about $_\oplus$ of iron while on the main sequence."564 They raise (he possibility (hat the Sun may have accreted a similar amount of iron during its evolution., They raise the possibility that the Sun may have accreted a similar amount of iron during its evolution.565 Whether helioseismologv can detect. this amount of accretion is uncertain., Whether helioseismology can detect this amount of accretion is uncertain.566 So far only one study has been published on the effect of accretion on the p-mode frequency (Henney Ulrich 1993)., So far only one study has been published on the effect of accretion on the -mode frequency (Henney Ulrich 1998).567 These authors concluded Chat the accretion of of meteoric malerial (approximately of iron) on the Sun could not be detected by seismology due io other uncertainties in the moclels., These authors concluded that the accretion of $_\oplus$ of meteoric material (approximately $_\oplus$ of iron) on the Sun could not be detected by seismology due to other uncertainties in the models.568 The main consequence of such accretion is (he metal enrichment of (he Sun's surface and convection zone. while maintaining a lower-abundance core.," The main consequence of such accretion is the metal enrichment of the Sun's surface and convection zone, while maintaining a lower-abundance core."569 In such a scenario. the expected neutrino [lux is reduced due to the lower opacity and temperature in (he central region of the Sun.," In such a scenario, the expected neutrino flux is reduced due to the lower opacity and temperature in the central region of the Sun."570 It has long been known that solar models with near-zero metallicity cores can lower the predicted neutrino flux. aud (his explanation was discussed earlv-on as a possible solution of (he classical neutrino problem (see e.g. the review by Rood 1973).," It has long been known that solar models with near-zero metallicity cores can lower the predicted neutrino flux, and this explanation was discussed early-on as a possible solution of the classical neutrino problem (see e.g. the review by Rood 1978)."571 More recently. Guenther Demarque (1997. hereafter GD97) have constructed solar models with low-Z cores using present physics ancl found their p-anode Irequencies to be incompatible with solar p-mocle observations.," More recently, Guenther Demarque (1997, hereafter GD97) have constructed solar models with low-Z cores using present physics and found their -mode frequencies to be incompatible with solar -mode observations."572 This paper describes work which can be viewed as an extension of GD97: but this lime we examine (he limits set by helioseismology on the non-standard assumption that the solar convection zone of the Sun has been enriched in heavy. elements by accretion during ils early evolution., This paper describes work which can be viewed as an extension of GD97; but this time we examine the limits set by helioseismology on the non-standard assumption that the solar convection zone of the Sun has been enriched in heavy elements by accretion during its early evolution.573 I addition to comparing the p-anode frequencies to observation in the frequency difference diagram. as done by CD97. we also compare the caleulated sound-speed and density in our models to (he same quantities derived Irom observation by inversion (Basu Antia 1997. hereafter DA97: Basu. Pinsonneault Baheall 2000. hereafter 11212001.," In addition to comparing the -mode frequencies to observation in the frequency difference diagram, as done by GD97, we also compare the calculated sound-speed and density in our models to the same quantities derived from observation by inversion (Basu Antia 1997, hereafter BA97; Basu, Pinsonneault Bahcall 2000, hereafter BPB00)."574 mince (his research was begun. the first results [rom the Sudbury Neutrino Observatory (SNO) experiment have been released. (Ahmad οἱ al.," Since this research was begun, the first results from the Sudbury Neutrino Observatory (SNO) experiment have been released (Ahmad et al."575 2001)., 2001).576 Ahmad et al. (, Ahmad et al. (577"2001) conclude that the D neutrino flux derived [rom their observations. ie. Px, = 5.44 4 0.99 x 10 ","2001) conclude that the $^{8}$ B neutrino flux derived from their observations, i.e. $\Phi_{^{8}B}$ = 5.44 $\pm$ 0.99 $\times$ $10^{6}$ "578uncerestimated errors in the 1990s. but in the last 12 vears jtve tracked Nz=1 very well.,"underestimated errors in the 1990s, but in the last 12 years have tracked $N_\sigma=1$ very well."579 The most recent 2 points for SN and 3 for BAO appear o have overestimated error bars. signficantlv so in the case of SN (bv a factor of 3).," The most recent 2 points for SN and 3 for BAO appear to have overestimated error bars, signficantly so in the case of SN (by a factor of 3)."580 SN measurements are those which most often quote systematic error bars (which we have added irectly to the statistical errors)., SN measurements are those which most often quote systematic error bars (which we have added directly to the statistical errors).581 We have tried two other wavs of dealing with the systematic errors. either. adding lem in quadrature. or ignoring them altogether.," We have tried two other ways of dealing with the systematic errors, either adding them in quadrature, or ignoring them altogether."582 We find that with the latter most conservative treatment. the SN results viel NV=0.5£0.07 and 0.77£0.12 for the most recent two points.," We find that with the latter most conservative treatment, the SN results yield $N_{\sigma}=0.5 \pm0.07$ and $0.77\pm0.12$ for the most recent two points."583 This is an improvement. indicating that the SN systematic error bars may well be too conservative.," This is an improvement, indicating that the SN systematic error bars may well be too conservative."584" Lt is still an underestimate. but now of similar magnitude to the cülferences seen between the accuracy-1 line aad some data points on the “other”. ""combined"" and “LSS” panels."," It is still an underestimate, but now of similar magnitude to the differences seen between the accuracy=1 line and some data points on the “other”, “combined” and “LSS” panels."585 Η we allow for the possibility that the Poisson error bars on our data points in Figure 13. are underestimates. and that there may be correlations between measurements in illerent. vears then this may σο some way to reconciling the measurements and their hoped for accuracy.," If we allow for the possibility that the Poisson error bars on our data points in Figure \ref{accm}586 are underestimates, and that there may be correlations between measurements in different years then this may go some way to reconciling the measurements and their hoped for accuracy."587 We return to lis point in our discussion below (Section 4.2)., We return to this point in our discussion below (Section 4.2).588 One question. which is not easy to answer from the multipanel Figures 12. and 13. is how the overall accuracy of measurements is changing by vear., One question which is not easy to answer from the multipanel Figures \ref{accp} and \ref{accm} is how the overall accuracy of measurements is changing by year.589 Are cosmological measurements improving as both theoretical knowledge and expertise in dealing with experimental uncertainties improve?, Are cosmological measurements improving as both theoretical knowledge and expertise in dealing with experimental uncertainties improve?590 We can see that this does appear to be the case by considering Figure 1H.. which plots accuracy by vear for results published. in the two main journals. ALINRAS and ApJ (including ApJL and ApJS).," We can see that this does appear to be the case by considering Figure \ref{journal}, which plots accuracy by year for results published in the two main journals, MNRAS and ApJ (including ApJL and ApJS)."591 These account for 354 ane 554 of all results in our compilation. respectively.," These account for $35\%$ and $55\%$ of all results in our compilation, respectively."592" The results before the vear ~2003 are significantly inaccurate. but steadily improve with time until after this date they become consistent. with the IN,=1 line."," The results before the year $\sim$ 2003 are significantly inaccurate, but steadily improve with time until after this date they become consistent with the $N_{\sigma}=1$ line."593 Both journals exhibit the same behaviour within their error bars., Both journals exhibit the same behaviour within their error bars.594 Using our tabulated data we can explore a few more aspects of the accuracy of measurements., Using our tabulated data we can explore a few more aspects of the accuracy of measurements.595 One can ask whether when results are published their accuracy alfects he amount by which they are cited. ancl therefore whether recognition increases with the perception that measurements are accurate.," One can ask whether when results are published their accuracy affects the amount by which they are cited, and therefore whether recognition increases with the perception that measurements are accurate."596 We address this in Figure. 15... where we Mot the accuracy vs. the number of citations to à paper. roth on a log scale.," We address this in Figure \ref{ncite}, where we plot the accuracy vs. the number of citations to a paper, both on a log scale."597 We can see that there appears. to » little evidence for any relationship between the two. so lat accuracy is not an important factor in determining re number of citations.," We can see that there appears to be little evidence for any relationship between the two, so that accuracy is not an important factor in determining the number of citations."598 Looking at Figure 15.. it does seem that there might. be slightly less papers with high Ny (innaccurate) and high. citations that other corners of 1e. plot.," Looking at Figure \ref{ncite}, , it does seem that there might be slightly less papers with high $N_{\sigma}$ (innaccurate) and high citations that other corners of the plot."599 This leads to a Pearson correlation coellicient of r—0.066. and therefore a slight. correlation between ótations ancl accuracy. in that papers with higher accuracy (lower Nz) have more citations.," This leads to a Pearson correlation coefficient of $r=-0.066$, and therefore a slight correlation between citations and accuracy, in that papers with higher accuracy (lower $N_{\sigma}$ ) have more citations."600 A set. of points with no correlation would give such a result 11% of the time. so the evidence for this is marginal. however.," A set of points with no correlation would give such a result $11\%$ of the time, so the evidence for this is marginal, however."601 We note that there does exist a significant. correlation between the precision of measurements and the number ofcitations (not plotted)., We note that there does exist a significant correlation between the precision of measurements and the number ofcitations (not plotted).602 We find a correlation coellicient. of, We find a correlation coefficient of603"burst of star formation. solar metallicity and using the 2. IME with lower and upper mass cutoffs ij,=0.1 M. and iii;=100 M...","burst of star formation, solar metallicity and using the \cite{Salpeter55} IMF with lower and upper mass cutoffs $m_L =0.1$ $_{\odot}$ and $m_U=100$ $_{\odot}$."604 Beddened model spectra. were created. [or the reddening range 0.0<E(5—V)«1.0 using the parametrized Galactic extinction law of ?.., Reddened model spectra were created for the reddening range $0.0 < E(B-V) < 1.0$ using the parametrized Galactic extinction law of \cite{Seaton79}.605 The five photometric UV/optical points suitable for the fit. i.e. those not affected by emission line contamination. were modelled using a minimum 4? technique to find the ages of the sec.," The five photometric UV/optical points suitable for the fit, i.e. those not affected by emission line contamination, were modelled using a minimum $\chi^{2}$ technique to find the ages of the SSC."606 Table 3 presents ranges of ages and recldenines For which acceptable fits were obtained (Nu 21): the ranges of measured masses are also presented.," Table 3 presents ranges of ages and reddenings for which acceptable fits were obtained $\chi^2_{red}607\lsim 1$ ); the ranges of measured masses are also presented."608 During the fitting process. the same (410%) error was assumed for all live photometric points.," During the fitting process, the same $\pm$ $\%$ ) error was assumed for all five photometric points."609 Taking in account the calibration error and the photon noise. both already discussed. in section 2.1. plus small shifts in the central wavelengths of the filters. and uncertainties due to the gradient in the background near the clusters. we believe (hat a mean error of d 1056 is realistic.," Taking in account the calibration error and the photon noise, both already discussed in section 2.1, plus small shifts in the central wavelengths of the filters, and uncertainties due to the gradient in the background near the clusters, we believe that a mean error of $\pm$ $\%$ is realistic."610 The best fitting models for the three clusters are shown in Figure 2., The best fitting models for the three clusters are shown in Figure 2.611 For comparison. some models (hat fall outside of the range of acceptable fits ave also shown.," For comparison, some models that fall outside of the range of acceptable fits are also shown."612 The latter represent the models with age/reddening combinations that give the best reduced chi-squared just outside the acceptable zone., The latter represent the models with age/reddening combinations that give the best reduced chi-squared just outside the acceptable zone.613 Taking the example of cluster C2 (see Figure 2b) the fit obtained Lor ihe 4 Myr with E(D-V)250.5 model is noticably better than that obtained for either the 3 Myr with E(B-Vj=0.6 or 6 Myr with E(D-V)20.4 models., Taking the example of cluster C2 (see Figure 2b) the fit obtained for the 4 Myr with E(B-V)=0.5 model is noticably better than that obtained for either the 3 Myr with E(B-V)=0.6 or 6 Myr with E(B-V)=0.4 models.614" Moreover 42, values found for YSP ages of 4-5 Mvr and reddening of E(B-V)=0.5 are. <1. while those found outside this range of ages ave 22."," Moreover $\chi^{2}_{red}$ values found for YSP ages of 4-5 Myr and reddening of E(B-V)=0.5 are $\lsim 1$, while those found outside this range of ages are $\gsim6152$."616 Therefore. (he range of ages considered to be valid for the case of C2 is 4-5 Myr. with the best fit found for a template of 4 Myr with a reddening of E(D-V)-0.5.," Therefore, the range of ages considered to be valid for the case of C2 is 4-5 Myr, with the best fit found for a template of 4 Myr with a reddening of E(B-V)=0.5."617 Note that the high fluxes observed for the FOC F320W point for the three clusters. rule out starburst ages above 50 Myr. even (aking in account any recldening effect. since the observed U—D is already bluer than such a starburst.," Note that the high fluxes observed for the FOC F320W point for the three clusters, rule out starburst ages above 50 Myr, even taking in account any reddening effect, since the observed $U-B$ is already bluer than such a starburst."618 Further evidence for the voung ages measured for the SSC in PINS1345+12 is provided by the detection of emission lines from the clusters in both the IST images (Figure 2) and the spectra (see next section)., Further evidence for the young ages measured for the SSC in PKS1345+12 is provided by the detection of emission lines from the clusters in both the HST images (Figure 2) and the spectra (see next section).619 Dased on the fluxes determined [rom our Ho. emission line images. and correcting for both the SSC continuum and a contribution from [NII[AA6548.6584. emission lines in the filter bandpass. we determine Πα equivalent widths of 250437. 370434 and 6062586. lor clusters C1. C2 and C4 respectively.," Based on the fluxes determined from our $\alpha$ emission line images, and correcting for both the SSC continuum and a contribution from $\lambda\lambda$ 6548,6584 emission lines in the filter bandpass, we determine $\alpha$ equivalent widths of $\pm$ 37, $\pm$ 34 and $\pm$ 86 for clusters C1, C2 and C4 respectively."620" The instantaneous burst models of ? show that such high equivalent widths are only attained lor clusters with ages /,4<G Myr. regardless of the assumed IMF."," The instantaneous burst models of \cite{Leitherer99}621 show that such high equivalent widths are only attained for clusters with ages $t_{cl} < 6$ Myr, regardless of the assumed IMF."622 Since (he ISM in the clusters may not absorb all the ionizine photons generated by the OB stars (as assumed by the Leitherer et al., Since the ISM in the clusters may not absorb all the ionizing photons generated by the OB stars (as assumed by the Leitherer et al.623 1999 models). the ages could," 1999 models), the ages could"624 (eg.Truupler Witt&Cordon1996.2000).. al.2006," \citep[e.g.,][]{trumpler30, mathis77, viallefond82, caplan85, witt92,625roussel05, driver08}. \citep[e.g.,][]{witt92, witt96, witt00}. \citep[e.g.,][]{waller92, deo06},"626).. 1996.2000). (6...Whittetetal.2001.2001).," \citep[e.g.,][]{elmegreen80, walterbos88, calzetti94,627witt96, witt00} \citep[e.g.,][]{vanhouten61, witt92, whittet01, whittet04}."628. observational scales. the dust will be intermixed with he stellar populations as thin lavers. fihuneuts. aud dense chimps.," observational scales, the dust will be intermixed with the stellar populations as thin layers, filaments, and dense clumps."629 Even though some lielt may be scattered into the Lue-ofsight from the rear and from reeious rear the line-ofsight. most of the extinction by far will ο due to dust that is distributed of the stellar xopulatious of interest (Byun1992).," Even though some light may be scattered into the line-of-sight from the rear and from regions near the line-of-sight, most of the extinction by far will be due to dust that is distributed of the stellar populations of interest \citep{byun92}."630". Popular inethods to measure dust extinction. such as he use of Uvdrogen recombination line ratios (e.g...Rudyuicuttctal. 2009).. the ultraviolet (UV) spectral slope (e.g...Calzettietal.1991:IKoug20010).. or ratios of he UV aud total infrared fluxes (οι,Buat&Xu1996:Calzettietal.2000:Boissier2001. 20053.. have various constraints and Πταος (οιοιNu&IInlouetal.2002:Doissier2001:Rieke 2001)."," Popular methods to measure dust extinction, such as the use of Hydrogen recombination line ratios \citep[e.g.,][]{rudy84, scoville01, calzetti05,631calzetti07, kennicutt09}, the ultraviolet (UV) spectral slope \citep[e.g.,][]{calzetti94, kong04}, or ratios of the UV and total infrared fluxes \citep[e.g.,][]{buat96, calzetti00, boissier04, boissier05}, have various constraints and limitations \citep[e.g.,][]{xu96, petersen97,632regan00, price02, boissier04, rieke04}."633.. Iu Tiuuuractal.(2009.hereafterPaperD.. we iutroduced a jew mcthod to measure the dust extinction in galaxies. and demoustrated that the spatial distribution of dust extinction in a late-tvpe spiral galaxy. 00959. can jo Inapped uxiug ouly theinages.," In \citet[][hereafter634Paper~I]{tamura09a}, we introduced a new method to measure the dust extinction in galaxies, and demonstrated that the spatial distribution of dust extinction in a late-type spiral galaxy, 0959, can be mapped using only the."635. Iu the preseut oper. we preseut the results of a iiulti:vaveleueth (UVopticaliud-IB) studw of the effect of the pixel-based wo-dimensional extinction correction of Paper I ou color conrposite images. a pixel Color Diagram (pCCD). and a pixel coordinate map of 00959.," In the present paper, we present the results of a multi-wavelength (UV--optical--mid-IR) study of the effect of the pixel-based two-dimensional extinction correction of Paper I on color composite images, a pixel Color–Color Diagram (pCCD), and a pixel coordinate map of 0959."636 These results show that a detailed pixel-based extinction correction Is essential to reveal the nature and distribution of stellar populations in galaxies., These results show that a detailed pixel-based extinction correction is essential to reveal the nature and distribution of stellar populations in galaxies.637turbulence when the region is dynamically bound €??)..,"turbulence when the region is dynamically bound \citep{Bateetal2003, Klessen2001}."638" On the small scales such as those represented by our protoclusters. Q.Ipe. the crossing time is of the order 10"" years."," On the small scales such as those represented by our protoclusters, 0.1pc, the crossing time is of the order $10^{5}$ years."639 We can therefore assume that all of our protoclusters form stars and that the star formation in our protoclusters takes place quickly. rapidly enough to be regarded here as instantaneous compared to the evolution of the whole GMC.," We can therefore assume that all of our protoclusters form stars and that the star formation in our protoclusters takes place quickly, rapidly enough to be regarded here as instantaneous compared to the evolution of the whole GMC."640 The simulation presented also has no method of incorporating feedback into the GMC model., The simulation presented also has no method of incorporating feedback into the GMC model.641 As is shown in figure in the right hand plot. the mass accreted into the gradually increases as the simulation progresses.," As is shown in figure in the right hand plot, the mass accreted into the gradually increases as the simulation progresses."642 At the point where the simulation is terminated. of the GMC has been accreted by the protoclusters.," At the point where the simulation is terminated, of the GMC has been accreted by the protoclusters."643 It is unlikely that this value is representative of how much mass would actually be involved in the star formation by this time. since feedback mechanisms such as ionisation. winds and supernovae would seriously alter the amount of gas that would be available for accretion into theSFCs.," It is unlikely that this value is representative of how much mass would actually be involved in the star formation by this time, since feedback mechanisms such as ionisation, winds and supernovae would seriously alter the amount of gas that would be available for accretion into the."644. What is needed is an estimate of when one would expect the star formation process to be halted by feedback mechanisms., What is needed is an estimate of when one would expect the star formation process to be halted by feedback mechanisms.645 This requires some knowledge of the star formation taking place within theSFCs., This requires some knowledge of the star formation taking place within the.646. We have already pointed out that the protoclusters in the simulation group into largeSFCs., We have already pointed out that the protoclusters in the simulation group into large.647. From now on in the paper we will use the details of these regions. rather than the individual »rotoclusters. to assess the nature of the star formation in the GMC.," From now on in the paper we will use the details of these regions, rather than the individual protoclusters, to assess the nature of the star formation in the GMC."648 Table gives the details of the after 9 Myr., Table gives the details of the after 9 Myr.649 The masses quoted for the in table includes all particles (SPH and ?rotoclusters) that fall within the radius of the region., The masses quoted for the in table includes all particles (SPH and protoclusters) that fall within the radius of the region.650 The gas yarticle Component is however quite small. generally less than20%.," The gas particle component is however quite small, generally less than."651 Although the star formation efficiency of GMCs is though ο be in the range of | toθέ. at the cluster level it is though ο be about 20 to depending on the region (for a discussion we point the reader to ?. and ?)).," Although the star formation efficiency of GMCs is thought to be in the range of 1 to, at the cluster level it is thought to be about 20 to depending on the region (for a discussion we point the reader to \citealt{Ladas2003} and \citealt{Kroupa2001}) )."652 In this paper we assume tha he star formation efficiency in our isSOG... but will include a discussion about the case in Wwich of the mass is turnec into stars.," In this paper we assume that the star formation efficiency in our is, but will include a discussion about the case in which of the mass is turned into stars."653 The assumed efficiency' here is high but this is deliberate since it actually assumes as little as possible about the effect tha the feedback mechanisms from t1e young stars are having on the accretion processes in theSFC., The assumed efficiency here is high but this is deliberate since it actually assumes as little as possible about the effect that the feedback mechanisms from the young stars are having on the accretion processes in the.654". We will also assume that the IMF of the stellar population in the follow a two step power law form. dNxm. ""dim. witha = S for LOS<afA.<0.5 and a Ξ 2.35 02). for 0.5«imfM.<100."," We will also assume that the IMF of the stellar population in the follow a two step power law form, $dN655\propto m^{-\alpha} dm$ , with $\alpha$ = 1.5 for $0.08 < m/M_{\odot} \le 0.5$ and $\alpha$ = 2.35 \citep{Salpeter1955} for $0.5 < m/M_{\odot} \le 100 $."656" This IMF. in conjunction with our assumption that o ""the mass of the is turned into stars. allows us to estimate tae Stellar population produced by the simulation."," This IMF, in conjunction with our assumption that of the mass of the is turned into stars, allows us to estimate the stellar population produced by the simulation."657 As already mentioned in the previous section. figure2 shows in the left hand plot how the mass of the 15 largest evolves with time.," As already mentioned in the previous section, figure shows in the left hand plot how the mass of the 15 largest evolves with time."658 The horizontal lines mark the point at which high mass stars can form., The horizontal lines mark the point at which high mass stars can form.659 From our IMF model. of the mass should be contained in stars Wwith masses greater thatLOM.," From our IMF model, of the mass should be contained in stars with masses greater that."660.. Thus a star will be present provided that there is 10/0.15 in the stellar populaticon., Thus a star will be present provided that there is 10/0.15 in the stellar population.661 Applying our assumed. star ormation efficiency of50%.. the must therefore have a mass of if they are to harbour a star.," Applying our assumed star formation efficiency of, the must therefore have a mass of if they are to harbour a star."662 The horizontal ong-dashed line in the figure denotes the point at which the achieve this mass., The horizontal long-dashed line in the figure denotes the point at which the achieve this mass.663 Doing the same for stars. which should comprise of the stellar mass in our chosen.. IMF. we tind that the need to contain 25/(0.0770.5) = if hey are to contain a star.," Doing the same for stars, which should comprise of the stellar mass in our chosen IMF, we find that the need to contain $25/(0.077664\times 0.5)$ = if they are to contain a star."665 This is represented by the yorizontal short-dashed line in the figure., This is represented by the horizontal short-dashed line in the figure.666 From our simple assumptions about the small scale efficiency and the form of the IMF. we ein estimate at what point in the simulation the star formation process will be disrupted by feedback mechanisms.," From our simple assumptions about the small scale efficiency and the form of the IMF, we can estimate at what point in the simulation the star formation process will be disrupted by feedback mechanisms."667 From figure2. we can estimate that the formation of stars would occur at about Cor at 4 Myr).," From figure, we can estimate that the formation of stars would occur at about (or at 4 Myr)."668 A star of mass would form after tor S Myr)., A star of mass would form after (or 5 Myr).669 Since the mass of the is increasing fairly rapidly at this point. stars with even higher masses would be expected to be present shortly after this. within 0.5 Myr or so.," Since the mass of the is increasing fairly rapidly at this point, stars with even higher masses would be expected to be present shortly after this, within 0.5 Myr or so."670 It would thus appear that the GMC is able to get enough mass into the for them to be able to form a full stellar population within about | Myr., It would thus appear that the GMC is able to get enough mass into the for them to be able to form a full stellar population within about 1 Myr.671 This is consistent with the observations of the small age spread in the stellar population ofthe Orion cluster (2)., This is consistent with the observations of the small age spread in the stellar population of the Orion cluster \citep{Hillenbrandetal2001}.672 Very rapidly after the tirst stars form we see that objects will be present., Very rapidly after the first stars form we see that objects will be present.673 This means that shortly after their formation. are going to contain ionising sources.," This means that shortly after their formation, are going to contain ionising sources."674 Such stars are commonly suggested to be responsible for controlling the star formation efficiency by expelling the gas from the cluster in which they form (such as our SFCs)). thus preventing the protostellar population from accreting or preventing new stars from forming.," Such stars are commonly suggested to be responsible for controlling the star formation efficiency by expelling the gas from the cluster in which they form (such as our ), thus preventing the protostellar population from accreting or preventing new stars from forming."675 However ? have noted that the ionisation from these stars does not appear to significantly affect the accretion rate in the clusters., However \citet{Daleetal2004} have noted that the ionisation from these stars does not appear to significantly affect the accretion rate in the clusters.676 The clumpy/fractal nature of the gas at the centre of the cluster where the OB type stars are situated acts to shield vast regions of the cluster from ionisation., The clumpy/fractal nature of the gas at the centre of the cluster where the OB type stars are situated acts to shield vast regions of the cluster from ionisation.677 Rather than pushing through the dense material. the ionising photons just find the path with the least resistance out of the cluster.," Rather than pushing through the dense material, the ionising photons just find the path with the least resistance out of the cluster."678 This is low density gas which would not normally be associated with protostellar accretion in the first place., This is low density gas which would not normally be associated with protostellar accretion in the first place.679 Similarly. the gas structure may also prevent the winds from OB stars expelling gas from the cluster.," Similarly, the gas structure may also prevent the winds from OB stars expelling gas from the cluster."680 It has been suggested that winds are able to escape via the fractal holes. without imparting much momentum to the dense regions (2)..," It has been suggested that winds are able to escape via the fractal holes, without imparting much momentum to the dense regions \citep{Henning1989}."681 It is therefore not clear if ionisation or winds will be able to expel the gas from cluster. thus halting the star formation process.," It is therefore not clear if ionisation or winds will be able to expel the gas from cluster, thus halting the star formation process."682 One mechanism that certainly will produce the desired effect is a supernova explosion., One mechanism that certainly will produce the desired effect is a supernova explosion.683 In fact it has been estimated that these events will not only remove the gas from a cluster. but also be able to disperse the natal GMC.," In fact it has been estimated that these events will not only remove the gas from a cluster, but also be able to disperse the natal GMC."684 Thus a high mass star's death will definitely mark the end of the star formation period in our cloud., Thus a high mass star's death will definitely mark the end of the star formation period in our cloud.685 Stars with masses greater than have very short sequence lifetimes. of about 3-5 Myr. and we see from the figure that they form at about 5 Myr after GMC formation.," Stars with masses greater than have very short main-sequence lifetimes, of about 3-5 Myr, and we see from the figure that they form at about 5 Myr after GMC formation."686 If we assume that a supernova event will occur at about 4+ Myr after the formation of the very high mass stars. then we estimate the first supernova event to occur at about 9 Myr or when the OB stars are 4 Myr old.," If we assume that a supernova event will occur at about 4 Myr after the formation of the very high mass stars, then we estimate the first supernova event to occur at about 9 Myr or when the OB stars are 4 Myr old."687 Assuming the supernova event will halt the star formation. we can now get an estimate of the star formation efficiency in the GMC.," Assuming the supernova event will halt the star formation, we can now get an estimate of the star formation efficiency in the GMC."688 The vertical dashed line in figure denotes the point at which we might see the first SN event., The vertical dashed line in figure denotes the point at which we might see the first SN event.689 At this time. 0.1 to 0.2 of the GMC's mass is contained in theSFCs.," At this time, 0.1 to 0.2 of the GMC's mass is contained in the."690. However we have assumed up until now that the efficiency in the is not butc.. therefore our estimate of the star formation efficiency in the GMC is roughly 5 to106c.," However we have assumed up until now that the efficiency in the is not but, therefore our estimate of the star formation efficiency in the GMC is roughly 5 to."691. This is easily comparable to the expected efficiencies in GMCs by Elmegreen's (2000) rapid cloud formation/dispersal model., This is easily comparable to the expected efficiencies in GMCs by Elmegreen's (2000) rapid cloud formation/dispersal model.692 The above analysis relied on a lot of assumptions about the nature of the star formation in theSFCs., The above analysis relied on a lot of assumptions about the nature of the star formation in the.693. In particular. it is guilty of invoking a star formation efficiency in the in order to determine the star formation efficiency. of the cloud: one could argue that this is not entirely self-consistent.," In particular, it is guilty of invoking a star formation efficiency in the in order to determine the star formation efficiency of the cloud: one could argue that this is not entirely self-consistent."694 Here we redo the above analysis butwithout the efficiency assumption., Here we redo the above analysis butwithout the efficiency assumption.695 If all the mass in the is used, If all the mass in the is used696be spherically svnimetric - see Oglev. Bell Burnell. Fender (2001) and Fender. Hanson. Pooley (1999).,"be spherically symmetric - see Ogley, Bell Burnell, Fender (2001) and Fender, Hanson, Pooley (1999)."697 We have observed the N-rav. binary Cvegnus X-3 with the ALERLIN interferometer with the intention of mapping jet components in à period of minor Dare activity., We have observed the X-ray binary Cygnus X-3 with the MERLIN interferometer with the intention of mapping jet components in a period of minor flare activity.698 We observed [or six epochs between OL December 1996 and 11 January 1907., We observed for six epochs between 01 December 1996 and 11 January 1997.699 In all epochs the source showed some degree of variability. ancl consequently. we were unable to map the source without the creation of artifacts., In all epochs the source showed some degree of variability and consequently we were unable to map the source without the creation of artifacts.700 Ligh time-resolution photometric observations show a number of small-Hux flares of around an hour in curation. with a large amount of structure.," High time-resolution photometric observations show a number of small-flux flares of around an hour in duration, with a large amount of structure."701 This is tvpical of a minor- period., This is typical of a minor-flare period.702 Power spectra of the data show no persistent »eriodicities in the data. and no evidence of the orbital »eriod of 4.8 hr.," Power spectra of the data show no persistent periodicities in the data, and no evidence of the orbital period of 4.8 hr."703 A measurement of the brightness temperature. for he Dares show typical values of 107-107. IX. with the argest values occurring over 4 minutes and a brightness emperature of 2«10H. qx. ‘Phe Hare. emission is [from a region of typically LO light minutes across. with a diameter of 1.22 AU.," A measurement of the brightness temperature for the flares show typical values of $10^{9}$ $10^{10}$ K, with the largest values occurring over 4 minutes and a brightness temperature of $2 \times70410^{11}$ K. The flare emission is from a region of typically 10 light minutes across, with a diameter of 1.22 AU."705 However. to be visible at a frequency of 5 Cllz hese plasmons would. be situated at a distance of 13 AU rom the core. assuming a spherical wind.," However, to be visible at a frequency of 5 GHz these plasmons would be situated at a distance of 13 AU from the core, assuming a spherical wind."706 The Authors wish to thank various MIZRLIN stall for help with cata and analysis., The Authors wish to thank various MERLIN staff for help with data and analysis.707 We are grateful for the help given by ‘Tom Muxlow. Simon Carrington and Peter Thomasson for scheduling the observations and. giving advice in the data reduction.," We are grateful for the help given by Tom Muxlow, Simon Garrington and Peter Thomasson for scheduling the observations and giving advice in the data reduction."708 The authors also wish to thank Duncan. Law-Green for his assistance in attempting to map he data., The authors also wish to thank Duncan Law-Green for his assistance in attempting to map the data.709 RNO wishes to hank Patrick MeCGrough and Robin Sanderson for their rospitality during the writing of this paper., RNO wishes to thank Patrick McGrough and Robin Sanderson for their hospitality during the writing of this paper.710 Guy Pooley generously made available the Itvle Tolescope data., Guy Pooley generously made available the Ryle Telescope data.711 The Green Bank Interferometer is a [facility of. the National Science. Foundation. operated. by he NILAO in support of NASA Lligh Inergv Astrophysics programme., The Green Bank Interferometer is a facility of the National Science Foundation operated by the NRAO in support of NASA High Energy Astrophysics programme.712 ALERLIN is a National Facility operated by the University of Manchester on behalf of PPARC., MERLIN is a National Facility operated by the University of Manchester on behalf of PPARC.713from Buttiglioneetal.(2010).,from \citet{but10}.714". The two main spectroscopic classes of Low and High Excitation Galaxies (LEG and HEG, respectively) define two separate sequences, which are also present in the SDSS emission line galaxies from Kewleyet (2006)."," The two main spectroscopic classes of Low and High Excitation Galaxies (LEG and HEG, respectively) define two separate sequences, which are also present in the SDSS emission line galaxies from \citet{kew06}."715". BL Lacertae falls in a region not well populated by 3CR sources, but it is closer to LEG than to HEG; furthermore it lies on the branch of the LINERs from the SDSS."," BL Lacertae falls in a region not well populated by 3CR sources, but it is closer to LEG than to HEG; furthermore it lies on the branch of the LINERs from the SDSS."716 This suggests a tentative identification as a Low Excitation Galaxy from the point of view of its narrow emission line spectrum., This suggests a tentative identification as a Low Excitation Galaxy from the point of view of its narrow emission line spectrum.717 Which class of objects share the properties of BL Lacertae described above?, Which class of objects share the properties of BL Lacertae described above?718" Let us start considering the radio-galaxies in the 3CR sample, see Fig. 5.."," Let us start considering the radio-galaxies in the 3CR sample, see Fig. \ref{3c}."719" The extended radio luminosity of BL Lacertae is at the faint end of the FR I in the 3CR, with only the nearby source 3C 272.1 (M 84) being fainter."," The extended radio luminosity of BL Lacertae is at the faint end of the FR I in the 3CR, with only the nearby source 3C 272.1 (M 84) being fainter."720" From the point of view of the narrow emission lines, the [OI] luminosity of BL Lacertae lies instead at the high end of 3CR/FR I. Thus this source is a strong outlier (by a factor of ~ 200—400 depending on the adopted value for its radio luminosity) from the relationship between line and radio luminosities followed by 3CRradio-galaxies?.."," From the point of view of the narrow emission lines, the ] luminosity of BL Lacertae lies instead at the high end of 3CR/FR I. Thus this source is a strong outlier (by a factor of $\sim$ 200–400 depending on the adopted value for its radio luminosity) from the relationship between line and radio luminosities followed by 3CR."721" This is true not only considering FR I, but also from a comparison of the LEG and HEG classes, and casts substantial doubts on the association of BL Lacertae with FR I, but more in general with the radio-galaxies in the 3CR sample."," This is true not only considering FR I, but also from a comparison of the LEG and HEG classes, and casts substantial doubts on the association of BL Lacertae with FR I, but more in general with the radio-galaxies in the 3CR sample."722" Note that also other classical samples of radio galaxies, like the B2 and the 2 Jy (Morgantietal.1997;Tadhunter1998) follow relations between radio and line emission similar to the 3CR."," Note that also other classical samples of radio galaxies, like the B2 and the 2 Jy \citep{morganti97,tadhunter98}723 follow relations between radio and line emission similar to the 3CR."724 Baldi&Capetti(2009) showed that a large ratio of line emission to the radio power with respect to 3CR sources is characteristic of the radio-loud AGN selected by Bestetal. (2005).., \citet{bal09} showed that a large ratio of line emission to the radio power with respect to 3CR sources is characteristic of the radio-loud AGN selected by \citet{best05a}.725" The latter authors cross-matched the ~212000 galaxies drawn from the SDSS-DR2 with the NVSS and radio surveys, selecting a sample of 2215 radio-loud AGN (with a radio flux threshold of 5 mJy) to which we refer hereafter as the SDSS/NVSS sample."," The latter authors cross-matched the $\sim726 212000$ galaxies drawn from the SDSS-DR2 with the NVSS and radio surveys, selecting a sample of 2215 radio-loud AGN (with a radio flux threshold of 5 mJy) to which we refer hereafter as the SDSS/NVSS sample."727 Baldi&Capetti(2010) explored the spectro-photometric properties of the SDSS/NVSS objects showing that they are generally hosted by massive early-type galaxies with a low excitation emission line spectrum.," \citet{bal10}728 explored the spectro-photometric properties of the SDSS/NVSS objects showing that they are generally hosted by massive early-type galaxies with a low excitation emission line spectrum."729" From the point of view of its narrow line luminosity and extended radio power BL Lacertae falls well within the region covered by the SDSS/NVSS sample (see ref3c)), although at a slightly higher L[OIIT] than the region of higher galaxies density."," From the point of view of its narrow line luminosity and extended radio power BL Lacertae falls well within the region covered by the SDSS/NVSS sample (see \\ref{3c}) ), although at a slightly higher $L \rm [\ion{O}{III}]$ than the region of higher galaxies density."730" However, these SDSS/NVSS sources do not show prominent broad lines like those observed in BL Lacertae."," However, these SDSS/NVSS sources do not show prominent broad lines like those observed in BL Lacertae."731" Because from the point of view of its radio flux “misaligned BL Lacertae"" objects would be included in the SDSS/NVSS catalogue even if located up to a redshift of ~0.2, the only possibility before we conclude that BL Lacertae is a unique object, at least in the nearby Universe, is that “misaligned BL Lacertae"" objects were rejected on an optical basis."," Because from the point of view of its radio flux “misaligned BL Lacertae"" objects would be included in the SDSS/NVSS catalogue even if located up to a redshift of $\sim 0.2$, the only possibility before we conclude that BL Lacertae is a unique object, at least in the nearby Universe, is that “misaligned BL Lacertae"" objects were rejected on an optical basis."732The recent discovery. of large numbers of. racdio-selected CUASIUS with very redo opticaltonear-infrared colours SUBEECSES that existing quasar surveys may be severely incomplete.,The recent discovery of large numbers of radio-selected quasars with very red optical–to–near-infrared colours suggests that existing quasar surveys may be severely incomplete.733 Based on a high identification rate in the optical and near-infrarecd.. Webster ct al. (," Based on a high identification rate in the optical and near-infrared, Webster et al. ("7341995). found. à. broad range of colours with 2«οfy<10 for Ilat-spectrum racio quasars in a subsample of the Parkes διτα] survey (Drinkwater ct al.,1995) found a broad range of colours with $2<B-K<10$ for flat-spectrum radio quasars in a subsample of the Parkes 2.7GHz survey (Drinkwater et al.735 1997: jereafter “Parkes quasars’)., 1997; hereafter `Parkes quasars').736 Lor comparison. quasars selected by standard optical techniques show a small scatter around οA~2.5.," For comparison, quasars selected by standard optical techniques show a small scatter around $B-K\sim2.5$."737 Two theories have »en proposed to explain the large scatter: first. Webster et al. (," Two theories have been proposed to explain the large scatter: first, Webster et al. ("7381995) interpreted. this in terms of extinction bv. line-ol-sight cust (see also Masci 1997). and second. Serjeant tawlines (1996) suggested that this was due to intrinsically red. optical/infrared svnchrotron radiation associated. with he radio emission.,"1995) interpreted this in terms of extinction by line-of-sight dust (see also Masci 1997), and second, Serjeant Rawlings (1996) suggested that this was due to intrinsically red optical/infrared synchrotron radiation associated with the radio emission."739 This paper explores a third. possibility: hat he red. colours are due to contamination by starlight rom the host. galaxies., This paper explores a third possibility: that the red colours are due to contamination by starlight from the host galaxies.740 In a recent near-infrarecl imaging study of a sample of raclio-quasars selected [rom the Iow-frequeney. (4308MIEIZ) catalog. Benn et al. (," In a recent near-infrared imaging study of a sample of radio-quasars selected from the low-frequency (408MHz) catalog, Benn et al. ("7411998). found that sources with red BOW colours to 2~ could be explained by an excess of host galaxy light in A.,1998) found that sources with red $B-K$ colours to $z\sim2$ could be explained by an excess of host galaxy light in $K$.742 Most. of the images appeared non-stellar (or resolved) suggesting that indeed starlight was responsible for the redness in 2A colour., Most of the images appeared non-stellar (or resolved) suggesting that indeed starlight was responsible for the redness in $B-K$ colour.743 I is importan to note however that all these sources were associated with extended radio galaxies of which a majority are of the steep-spectrum tvpe. à common feature of low-frequency: selectec samples.," It is important to note however that all these sources were associated with extended radio galaxies of which a majority are of the steep-spectrum type, a common feature of low-frequency selected samples."744 A Large [fraction are also often associated. with luminous ὃνgiant ellipticals., A large fraction are also often associated with luminous giant ellipticals.745 In view of the simple orientation-based. unified model for racio-louc AGN. it is possible tha ealaxy light is more easily detected in these sources due to anisotropic obscuration and/or unbeamed. emission. of the central AGN.," In view of the simple orientation-based unified model for radio-loud AGN, it is possible that galaxy light is more easily detected in these sources due to anisotropic obscuration and/or unbeamed emission of the central AGN."746 The radio-quasars explored in this paper are all of the fIat-spectrum type selected at moderately high requeney., The radio-quasars explored in this paper are all of the flat-spectrum type selected at moderately high frequency.747 Εις. they are not expected to exhibit. similar woperties in the near-inlrarect.," Thus, they are not expected to exhibit similar properties in the near-infrared."748 The 5ἐν colours of normal radio galaxies are known o be quite red. exhibiting a similar dispersion to those observed. in the Parkes sample.," The $B-K$ colours of normal radio galaxies are known to be quite red, exhibiting a similar dispersion to those observed in the Parkes sample."749 These. sources are. often associated with eiant ellipticals ancl their colours appear o be uniformly distributed. over the range 35D——NET or redshifts 22x2 (MeCarthy1993andreferencestherein)., These sources are often associated with giant ellipticals and their colours appear to be uniformly distributed over the range $3\simlt B-K\simlt 7$ for redshifts $z\simlt2$ \cite{McCarthy1993}.750.. Lt is possible that the host galaxies of Parkes quasars also exhibit similar properties., It is possible that the host galaxies of Parkes quasars also exhibit similar properties.751 To determine the importance of vost galaxy Light in Parkes quasars. we need to quantify its contribution throughout the optical to near-Ht wavelength reeion.," To determine the importance of host galaxy light in Parkes quasars, we need to quantify its contribution throughout the optical to near-IR wavelength region."752 Determining the host galaxy properties of radio-cuasars and DL-Lac-twvpe sources are crucial for studies of ACN evolution and testing unified schemes., Determining the host galaxy properties of radio-quasars and BL-Lac-type sources are crucial for studies of AGN evolution and testing unified schemes.753 A significant number of BL-Lacs and other core-dominated racio Lou: AGN are, A significant number of BL-Lacs and other core-dominated radio loud AGN are754X-ray. transient. sources are X-ray. binary. svstenmis. typically. with a low mass companion star and a black hole for the compact object (although sometimes a neutron star).,"X-ray transient sources are X-ray binary systems, typically with a low mass companion star and a black hole for the compact object (although sometimes a neutron star)."755 They are well-known for dramatic outbursts caused by some form instability ↔within.. the. accretion Dow., They are well-known for dramatic outbursts caused by some form of instability within the accretion flow.756s There ofare now  40> sources. some of which (c.g.are recurrent although the ⋅ majority have still only been observed in outburst once.," There are now $\sim 40$ observed sources, some of which are recurrent although the majority have still only been observed in outburst once."757" The canonical “soft” X-ray. transient. outburst is one which displavs a soft. blackbody spectrum in low energy X-ravs. and a “Fast. Rise"" Exponential:. Decay” oq.lighteurve morphology (e.g. Chen.. Shracer. LivioD 1997)."," The canonical “soft” X-ray transient outburst is one which displays a soft blackbody spectrum in low energy X-rays and a “Fast Rise Exponential Decay” lightcurve morphology (e.g. Chen, Shrader Livio 1997)."758B WithL modern X-ray telescopes it has been possible to obtain the spectral and temporal coverage to show that the behaviour is. more complicated. than this., With modern X-ray telescopes it has been possible to obtain the spectral and temporal coverage to show that the behaviour is more complicated than this.759 ⋠↽X-ray transients. appear to enter the outburst from an initially hard spectral state (the low/hard state: see e.g. van der Klis 1995)., X-ray transients appear to enter the outburst from an initially hard spectral state (the low/hard state; see e.g. van der Klis 1995).760 On a timescale ~⋅ days. weeksthe X-ray source⋅ then softens in⊀ most⋅⊀ cases of Brocksopp etonnoy al., On a timescale of $\sim$ days–weeks the X-ray source then softens in most cases (e.g. Brocksopp et al.761 2002)., 2002).762 @Some X-ray. transients.. however.⋅⋠ observed not soften but⊀⋠ remain in dothe Iow/hard state throughout the outburst (e.g. Brocksopp et al.," Some X-ray transients, however, do not soften but remain in the low/hard state throughout the outburst (e.g. Brocksopp et al."763 2004. 2001).," 2004, 2001)."764 All black hole X-ray transient sources which have been observed. at radio ⊀⋅frequencies. have been detected. at some point. during. their. lon(Brocksopp et al., All black hole X-ray transient sources which have been observed at radio frequencies have been detected at some point during their (Brocksopp et al.765 nno2002).," 2002),"766COCCHELICIEY resonance Is as ollows: on the nominal resonance. the two planets have successive conjunctions at exactly. the same longitude in inertial soie.,"eccentricity resonance is as follows: on the nominal resonance, the two planets have successive conjunctions at exactly the same longitude in inertial space."767 The strong interactions that occur at conjunctions build up the cecentricity of the test xuwticle and cause a change in semimajor axis ancl period., The strong interactions that occur at conjunctions build up the eccentricity of the test particle and cause a change in semimajor axis and period.768 The change in period of the test. particle causes the longitude of conjunction to drift., The change in period of the test particle causes the longitude of conjunction to drift.769 Once 1e longitude of conjunction shifts by about x relative to the original direction. the eccentricity »eeins to decrease making a libration evcle.," Once the longitude of conjunction shifts by about $\pi$ relative to the original direction, the eccentricity begins to decrease making a libration cycle."770 The libration of the semi-major axes causes the timing of the transits to change., The libration of the semi-major axes causes the timing of the transits to change.771 This qualitative discussion leads directly to an estimate of the drifts in transit times., This qualitative discussion leads directly to an estimate of the drifts in transit times.772 Within each libration evcle the ongitude of conjunction shifts by about half an orbit. mostly due to the period change of the lighter planet.," Within each libration cycle the longitude of conjunction shifts by about half an orbit, mostly due to the period change of the lighter planet."773 Since conjunctions occur only once every j ortnits the largest transit time deviation of the lighter planet during the period of libration is £j (in this order of magnitucle derivation we ignore factors of order unity. and take the limit of large j so that j2|1 and Pox DU.," Since conjunctions occur only once every $j$ orbits the largest transit time deviation of the lighter planet during the period of libration is $P/j$ (in this order of magnitude derivation we ignore factors of order unity, and take the limit of large $j$ so that $j\simeq j+1$ and $P_2\simeq P_1$ )."774 The observationavy more interesting case is probably that in which the heavier planet is the transiting one., The observationally more interesting case is probably that in which the heavier planet is the transiting one.775 Then. conservation of energy. [or 1¢ orbiting planets implies that the change in periods is inversely. proportional to the masses. herefore the timing variations are given by (miniémis)£j.," Then, conservation of energy for the orbiting planets implies that the change in periods is inversely proportional to the masses, therefore the timing variations are given by $(m_{light}/m_{heavy})P/j$."776 Wo lind an excellent fit to the data for The caleulations shown in Figure 3.2 verily this analytic scaling with 7., We find an excellent fit to the data for The calculations shown in Figure \ref{fig3} verify this analytic scaling with $j$.777 Calculating the libration period is a little more complicated. but still straightforward.," Calculating the libration period is a little more complicated, but still straightforward."778 Suppose the period of the test yarticle deviates from the nominal resonance by a small fraction e., Suppose the period of the test particle deviates from the nominal resonance by a small fraction $\epsilon$.779 Phen. consecutive conjunctions drift in longitude by about E c.," Then, consecutive conjunctions drift in longitude by about $2\pi j^2 \epsilon$ ."780" The number of conjunctions. IN... before a drift of order x in the longitude of conjunctions accumulates is IN,72τεl"," The number of conjunctions, $N_c$, before a drift of order $\pi$ in the longitude of conjunctions accumulates is $N_c \sim j^{-2}\epsilon^{-1}$."781 We now estimate e indirectly., We now estimate $\epsilon$ indirectly.782 Phe test particle gains an eccentricity of order j7 fiin each conjunction due to the radial force from he massive planet (this can be computed from the impulse approximation and the perturbation equation for eccentricity)., The test particle gains an eccentricity of order $j^2\mu$ in each conjunction due to the radial force from the massive planet (this can be computed from the impulse approximation and the perturbation equation for eccentricity).783 The cecentricity given in Ny conjunctions is then of order Ae~ pret., The eccentricity given in $N_c$ conjunctions is then of order $\Delta e \sim \mu \epsilon^{-1}$ .784 Using the Tisserand relation. the fractional change in semimajor. axis. associated. with. this. change in. eccentricityma is. joe⋅27.," Using the Tisserand relation, the fractional change in semimajor axis associated with this change in eccentricity is $j \mu^2 \epsilon^{-2}$."785 Sincenm this. is. also the fractionalH. change in. the period.. we rave cEIT7 and a libration period. of We numerically computed the amplitude and period of the transit timing variations at the 2:1 resonance., Since this is also the fractional change in the period we have $\epsilon \sim j^{1/3} \mu^{2/3}$ and a libration period of We numerically computed the amplitude and period of the transit timing variations at the 2:1 resonance.786 Figure 4. shows a plot of the amplitude of the timing variations versus the mass ratio of the perturbing planet to the transiting planet., Figure \ref{fig4} shows a plot of the amplitude of the timing variations versus the mass ratio of the perturbing planet to the transiting planet.787 As »edieted. the amplitude is of order the period of the transiting planet when the transiting planet is lighter. ancl varies as the mass ratio when the transiting planet is heavier.," As predicted, the amplitude is of order the period of the transiting planet when the transiting planet is lighter, and varies as the mass ratio when the transiting planet is heavier."788 The libration period measured from the numerical simulations shows the xedieted behavior. scaling precisely as ye777 for the more massive planet (with a coelficient of ~0.7 for j=1 and 0.5 for jc ldn equation 34)).," The libration period measured from the numerical simulations shows the predicted behavior, scaling precisely as $\mu^{-2/3}$ for the more massive planet (with a coefficient of $\sim 0.7$ for $j=1$ and $0.5$ for $j>1$ in equation \ref{plib}) )."789" We have compared the numerical values of the amplitude and period of libration on resonance as a ""unction of j.", We have compared the numerical values of the amplitude and period of libration on resonance as a function of $j$.790 Despite the fact that the above scalings were derived in the large-j limit. the agreement is better than 10 per cent lor jzc2. and accurate to about 40 per cent for j=1.," Despite the fact that the above scalings were derived in the $j$ limit, the agreement is better than 10 per cent for $j\ge 2$, and accurate to about 40 per cent for $j=1$."791 Figure 4 shows the more detailed behavior of the amplitude near the 2:1 resonance., Figure \ref{fig4} shows the more detailed behavior of the amplitude near the 2:1 resonance.792 The amplitude is maximum slightIv »dlow resonance at the location of the cusp., The amplitude is maximum slightly below resonance at the location of the cusp.793 This may be understood. as follows: since the simulations are started. with CL—eo=0. after conjunction the eccentricity grows and the outer planet moves outwards. while the inner planet. moves inward.," This may be understood as follows: since the simulations are started with $e_1=e_2=0$, after conjunction the eccentricity grows and the outer planet moves outwards, while the inner planet moves inward."794 This causes the planets to move closer to resonance. causing a longer time between conjunctions. leading to a larger change in eccentricity and semi-major axis.," This causes the planets to move closer to resonance, causing a longer time between conjunctions, leading to a larger change in eccentricity and semi-major axis."795 The cusp is the location where the planets reach exact. resonance at the turning point of libration. at which point 6 is maximum.," The cusp is the location where the planets reach exact resonance at the turning point of libration, at which point $\delta t$ is maximum."796 To the right of the cusp. the libration causes the planets to overshoot the resonance. so the change in eccentricity ancl semi-major axis is somewhat smaller. and hence the amplitude is smaller.," To the right of the cusp, the libration causes the planets to overshoot the resonance, so the change in eccentricity and semi-major axis is somewhat smaller, and hence the amplitude is smaller."797 ligure 4. shows that the width of the resonance scales as στ)5 (the horizontal axis has been scaled with ji2* so that the curves overlap). so for larger mass planets the resonant variations have a wider range of influence than the non-resonant variations discussed in the previous section.," Figure \ref{fig4} shows that the width of the resonance scales as $\mu^{2/3}$ (the horizontal axis has been scaled with $\mu^{-2/3}$ so that the curves overlap), so for larger mass planets the resonant variations have a wider range of influence than the non-resonant variations discussed in the previous section."798 The curves in Figure 4. demonstrate that on-resonance the amplitude scales as mintl.pposunffuraos)/j- while olf-resonance the amplitude scales as pipers.," The curves in Figure \ref{fig4}799 demonstrate that on-resonance the amplitude scales as $min(1,\mu_{pert}/800\mu_{trans})/j$, while off-resonance the amplitude scales as $\mu_{pert}$."801 When either eccentricity is [large enough. higher order resonances become important.," When either eccentricity is large enough, higher order resonances become important."802 In. particular. the resonances that are lim begin o dominate as the ratio of the semi-major axes becomes large: as the cecentricity of the outer planet approaches unity these resonances become as strong as first order resonances (2)..," In particular, the resonances that are $m$ begin to dominate as the ratio of the semi-major axes becomes large; as the eccentricity of the outer planet approaches unity these resonances become as strong as first order resonances \citep{pan04}."803 LFigure 5. shows the results of a numerical calculation where the ransiting planet Η 200458» with a mass of approximately 0.67. Jupiter masses. is perturbed by a TAL) planet with various eccentricities (we have taken LID 209458b to have à circular orbit).," Figure \ref{fig5} shows the results of a numerical calculation where the transiting planet HD 209458b, with a mass of approximately 0.67 Jupiter masses, is perturbed by a $1{\rm M_\oplus}$ planet with various eccentricities (we have taken HD 209458b to have a circular orbit)."804 Near the mean-motion resonances the signal is large enough that an carth-mass planet would be detectable with current technology., Near the mean-motion resonances the signal is large enough that an earth-mass planet would be detectable with current technology.805 “Phe amplitude increases everywhere with eccenricitv., The amplitude increases everywhere with eccentricity.806 This graph can be applied to systems with other massesand periods as the timing variationscales as Πλ.Mrere Cexcopt for planets trapped in resonance), This graph can be applied to systems with other massesand periods as the timing variationscales as $\delta t \propto P_{trans} m_{pert}$ (except for planets trapped in resonance).807 When roth planets have non-zero eccentricity. the parameter space becomes quite large: the + phase space coordinates for," When both planets have non-zero eccentricity, the parameter space becomes quite large: the 4 phase space coordinates for"808 , 809us more detailed information. however. a DEM analysis was not possible with our data set because of the small number of emission lines available.,"us more detailed information, however, a DEM analysis was not possible with our data set because of the small number of emission lines available."810 Assuming an isothermal plasma. the radiative loss from microflares is estimated by using a total radiative loss function.," Assuming an isothermal plasma, the radiative loss from microflares is estimated by using a total radiative loss function."811 Since the emission measure of the microflare peaks around 1x10° K as discussed above. the column emission measure at that temperature is derived from Equation 1. by using the 419.512 nm radiance at the peak time of the lighteurves.," Since the emission measure of the microflare peaks around $1\times10^6$ K as discussed above, the column emission measure at that temperature is derived from Equation \ref{eq:em}812 by using the $\lambda$ 19.512 nm radiance at the peak time of the lightcurves."813 The column emission measure is multiplied by the apparent area of the microflare to obtain a total emission measure., The column emission measure is multiplied by the apparent area of the microflare to obtain a total emission measure.814 A radiative loss function is derived from the CHIANTI atomic database assuming a typical density of 3x10° em? from Table 2.., A radiative loss function is derived from the CHIANTI atomic database assuming a typical density of $3\times10^9$ $^{-3}$ from Table \ref{table:bp}.815" The radiative loss function is estimated to be 4x107 erg eni? s! at the temperature of |x10° K. Radiative loss rates of microflares are calculated by multiplying the radiative loss function and the total emission measure, which are listed in the last column of Table 2.."," The radiative loss function is estimated to be $4\times10^{-22}$ erg $^3$ $^{-1}$ at the temperature of $1\times10^6$ K. Radiative loss rates of microflares are calculated by multiplying the radiative loss function and the total emission measure, which are listed in the last column of Table \ref{table:bp}."816 The radiative loss rates range from 2x107! to 6x107 erg s., The radiative loss rates range from $2\times10^{21}$ to $6\times10^{22}$ erg $^{-1}$.817 As shown in Figs., As shown in Figs.818 10 and 11.. actual microflares are not isothermal.," \ref{fig:loci01} and \ref{fig:loci08}, actual microflares are not isothermal."819" Therefore. our results are considered as the lower limit of the radiative loss rate since the remaining component could also contribute to the radiative loss,"," Therefore, our results are considered as the lower limit of the radiative loss rate since the remaining component could also contribute to the radiative loss."820 The observed delay times between the peaks in X-ray and cool emissions are in the 8 — 21 min range (Table 2))., The observed delay times between the peaks in X-ray and cool emissions are in the 8 – 21 min range (Table \ref{table:bp}) ).821 Theoretical cooling times are computed to determine if the delay times can be attributed to the cooling of hot coronal loops., Theoretical cooling times are computed to determine if the delay times can be attributed to the cooling of hot coronal loops.822 ? derived an analytical solution for the cooling time from a combination of conductive cooling and radiative cooling. taking into account chromospheric evaporation.," \citet{cargill1995} derived an analytical solution for the cooling time from a combination of conductive cooling and radiative cooling, taking into account chromospheric evaporation."823" It requires temperature Το. density Πο. and loop length {, as initial parameters (from Equation 14E "," It requires temperature $T_0$, density $n_e$, and loop length $L$ as initial parameters (from Equation 14E in \citealt{cargill1995}) )."824Although the internal structure. of microflares is. not resolved in our observations. the loop size 1s assumed to be," Although the internal structure of microflares is not resolved in our observations, the loop size is assumed to be"825The mean velocity for each cluster is listed in Table 10..,The mean velocity for each cluster is listed in Table \ref{samplevr}.826 Most of the values obtained agree. within (he uncertainties. wilh previous measurements from the literature. even in the case of the clusters observed with the VLT. where the uneertainties are larger.," Most of the values obtained agree, within the uncertainties, with previous measurements from the literature, even in the case of the clusters observed with the VLT, where the uncertainties are larger."827 In the case ol NGC 2141. we found a mean velocity similar to the value obtained by Coleetal.(2004).," In the case of NGC 2141, we found a mean velocity similar to the value obtained by \citet{c04}."828. Both values differ by 20 and 30 km !. respectively. from (he value found by (2002).," Both values differ by 20 and 30 km $^{-1}$, respectively, from the value found by \citet{f02}."829. For Collinder 110. no previous measurement of its radial velocity could be found in ihe literature.," For Collinder 110, no previous measurement of its radial velocity could be found in the literature."830 We are interested in obtaining metallicities Irom red giant stars. ancl within (his eroup. rom (he brightest ones. which are of spectral tvpes Ix ancl M. The main features in the infrared spectra of these stars are the CaT lines.," We are interested in obtaining metallicities from red giant stars, and within this group, from the brightest ones, which are of spectral types K and M. The main features in the infrared spectra of these stars are the CaT lines."831 But (heir spectra also contains other weak atomic lines., But their spectra also contains other weak atomic lines.832 The FeI (8514.1. 8674.8. 8688.6 and 8824.2 À)) and Ti (8435.0 À)) lines are the nost important.," The Fe (8514.1, 8674.8, 8688.6 and 8824.2 ) and Ti (8435.0 ) lines are the most important."833 When within this range. we move to later spectral types. and hence to cooler stars. molecular bands begin to appear that change the slope of the local continuum.," When within this range, we move to later spectral types, and hence to cooler stars, molecular bands begin to appear that change the slope of the local continuum."834 The nain contribution are from the titanium oxide (TiQ) bands. the strongest of which are the triplet situated at S432. 8442 and 8452 ancl the doublet at 8859.6 and 8863.5A.," The main contribution are from the titanium oxide (TiO) bands, the strongest of which are the triplet situated at 8432, 8442 and 8452 and the doublet at 8859.6 and 8868.5."835. There are other weaker bands at 8472. 8506. 8513. 8558 and 8569À.. near the bluest lines of the CaT. There are also several vanadium oxide (VO) bands at 8521. 8538. 8574. 8597. 8605. 8624. 8649 and 8668A.," There are other weaker bands at 8472, 8506, 8513, 8558 and 8569, near the bluest lines of the CaT. There are also several vanadium oxide (VO) bands at 8521, 8538, 8574, 8597, 8605, 8624, 8649 and 8668."836. The strength of these features increases when (he temperature decreases. Wwhen we move to later spectral (vpes.," The strength of these features increases when the temperature decreases, when we move to later spectral types."837 The presence of these bands complicates the definition of the continuum. which makes it difficult to obtain the equivalent widths of the CaT lines for stars with Ἑν 23500 IX or (V-I)22. in the most metal-rich clusters.," The presence of these bands complicates the definition of the continuum, which makes it difficult to obtain the equivalent widths of the CaT lines for stars with $_{eff}\leq$ 3500 K or $>$ 2, in the most metal-rich clusters."838 The description of the CaT reeion for other spectral tvpes can be found in Cenarroetal.(2001)., The description of the CaT region for other spectral types can be found in \citet{cen01}.839. In the literature we can lind different prescriptions to measure the strength of the CaT lines., In the literature we can find different prescriptions to measure the strength of the CaT lines.840 The classical definition of a spectral index consists in establishing a central bandpass covering a spectral feature and one or more bandpasses on both sides to trace the local continuum relerence level., The classical definition of a spectral index consists in establishing a central bandpass covering a spectral feature and one or more bandpasses on both sides to trace the local continuum reference level.841 Cenarroetal.(2001) have presented a description of the previous CaT index definitions and a comparison among them., \citet{cen01} have presented a description of the previous CaT index definitions and a comparison among them.842 In Figure 3. we have plotted the line and continuum bandpasses used in several reference works. Cenarroetal.(2001). (a). Rutlecdeeetal.(1997a) (b) and Armandrolf&Zinn(1933). (ο). over a metal-poor (left) and a metal-rich (right) spectrum.," In Figure \ref{bandas} we have plotted the line and continuum bandpasses used in several reference works, \citet{cen01} (a), \citet{r97a} (b) and \citet{az88} (c), over a metal-poor (left) and a metal-rich (right) spectrum."843 The Armancdroff&Zinn(1988) and Rutledgeetal.(19974) , The \citet{az88} and \citet{r97a} 844The general case with z€E! could be deduced from (2.21)) by trauslation.,The general case with $z\in \R^{n+1}$ could be deduced from \ref{kersha}) ) by translation.845 Throughout the prool. we will sometimes omit (when there is no coufusion) the cdepeucence of the uorm on the space E+.," Throughout the proof, we will sometimes omit (when there is no confusion) the dependence of the norm on the space $\R^{n+1}$."846 The proof is divided into three (Decomposition of (6;*w)(0) ou parabolic Since o; is supported in {2&E5;»5/τςzcoT then o;(0)20—fpo).," The proof is divided into three (Decomposition of $(\phi_{j}*u)(0)$ on parabolic Since $\h{\phi}_{j}$ is supported in $\{z\in \R^{n+1};\,8472^{j-1}\leq \|z\|\leq 2^{j+1}\}$ then $\h{\phi}_{j}(0)=0=\int_{\R^{n+1}}\phi_{j}$."848 Using this equality. we can write: where Qt/ is the parabolie cube defined by (2.20)) and centered at 0.," Using this equality, we can write: where $Q^{1-j}$ is the parabolic cube defined by \ref{key_para}) ) and centered at $0$."849" This implies that (Estimate of A4). From (2.7)). the term Ay can be estimated as follows: hence (Estimate of As). We rewrite A> as the following series: Since ©,1 is the inverseFourier transform of a compactly supported function then we have: The asvinptotic behavior of ó4 shown by (2.23)) leads to the followiug clecomposition of the term. Ap:"," This implies that (Estimate of $A_{1}$ From \ref{hedshsu1}) ), the term $A_{1}$ can be estimated as follows: hence (Estimate of $A_{2}$ We rewrite $A_{2}$ as the following series: Since $\phi_{1}$ is the inverseFourier transform of a compactly supported function then we have: The asymptotic behavior of $\phi_{1}$ shown by \ref{key_eq4}) ) leads to the following decomposition of the term $A_{2}$ :"850" For the five subsamples from SSLOT. core raclio luminosities (in the form £L, ,,) and the resulting radio Loudness paranicters (Lis,/ Lo) are presented in‘Tables Al to AX5.."," For the five subsamples from SSL07, core radio luminosities (in the form $\nu_{5} L_{\nu_{5, \rm \: core}}$ ) and the resulting radio loudness parameters $L_{\nu_{5, \rm \: core}}/L_{\nu_{B}}$ ) are presented inTables \ref{table: radio cores1} to \ref{table: radio cores5}. ."851Tex Collisional cascades plav a central role in planet formation.,7ex Collisional cascades play a central role in planet formation.852 Ii current theory. planets erow [rom collisions and mergers of km-sized plauetesimals in a gaseous disk.," In current theory, planets grow from collisions and mergers of km-sized planetesimals in a gaseous disk."853 As planets grow. they stir leftover planetesimals along their orbits to high velocities.," As planets grow, they stir leftover planetesimals along their orbits to high velocities."854 Eventually. collisions among planetesimals produce smaller fragments instead of larger. merged objects.," Eventually, collisions among planetesimals produce smaller fragments instead of larger, merged objects."855 Continued stirring leads to a cascade of destructive collisions which grinds the leftovers to dust., Continued stirring leads to a cascade of destructive collisions which grinds the leftovers to dust.856 This process (1) explains the masses of terrestrial planets (Ixenvon&Bromley2006) and Ixuiper belt objects (IXenvonοἱal.2008) and Gi) produces debris disks similar to those observed around nearby main sequence stars (Wyatt2003)., This process (i) explains the masses of terrestrial planets \citep{kb06} and Kuiper belt objects \citep{kbod08} and (ii) produces debris disks similar to those observed around nearby main sequence stars \citep{wya08}.857. Numerical simulations of icv planet lormation suggest the cascade limits the masses of erowing protoplanets to ~ 0.01 citep||[herealterINDOS| kbOs..," Numerical simulations of icy planet formation suggest the cascade limits the masses of growing protoplanets to $\sim$ 0.01 \\citep[][hereafter KB08]858{kb08}."859 ‘This mass is much smaller Chan (the core mass. = 0.11 required for a protoplanet to accrete gas and become a gas giant planet (Pollacketal.1996:|2003:Alibertetal.—. 2005).," This mass is much smaller than the core mass, $\gtrsim$ 0.1–1, required for a protoplanet to accrete gas and become a gas giant planet \citep{pol96,ina03,ali05}."860. Unless icy protoplanets can accrete collision fragments before the fragments ave ground to dust. these protoplanets cannot grow into gas ejant planet cores.," Unless icy protoplanets can accrete collision fragments before the fragments are ground to dust, these protoplanets cannot grow into gas giant planet cores."861 Thus. finding a mechanism (o halt the cascade is essential to form gas eijant planets.," Thus, finding a mechanism to halt the cascade is essential to form gas giant planets."862 llere. we describe how interactions between the fragments and the gaseous disk can halt the cascade.," Here, we describe how interactions between the fragments and the gaseous disk can halt the cascade."863 In our picture. (he gas (raps small lragments with sizes of 0.1 mun (o 1 m and prevents them [rom colliding at large velocities.," In our picture, the gas traps small fragments with sizes of 0.1 mm to 1 m and prevents them from colliding at large velocities."864 These Iragments then settle rapidly to the disk midplane. where protoplanets can accrete them.," These fragments then settle rapidly to the disk midplane, where protoplanets can accrete them."865 For a broad range of initial conditions. analvlic results and detailed nunerical simulations demonstrate that (his process vields ccores in 12 Myr.," For a broad range of initial conditions, analytic results and detailed numerical simulations demonstrate that this process yields 1--10 cores in 1–2 Myr."866function of distance.,function of distance.867 As the distance rises. d; gets larger and (1+2)πλ gets smaller. which gives a maximum value for the right side.," As the distance rises, $d_L^2$ gets larger and $(1+\textit{z})^{-3.04}$ gets smaller, which gives a maximum value for the right side."868 When the concordance cosmology is used. the funcion peaks al ο3.6.," When the concordance cosmology is used, the function peaks at $\textit{z}\sim3.6$."869 Specifically. the right side of the equation cannot exceed 1.13x10?keV?!ere!em? ancl. therefore. This becomes a simple way to test the Amati relation even for bursts without redshilts.," Specifically, the right side of the equation cannot exceed $1.13 \times 10^9 \unit{keV^{2.04} \; erg^{-1} \; cm^{2}}$ and, therefore, This becomes a simple way to test the Amati relation even for bursts without redshifts."870" Similarly. for the Ghirlanda relation. The beaming factor. £j,,,,. is defined as (1—cos0,,;). where 9;,, is the opening angle of the jet of the burst."," Similarly, for the Ghirlanda relation, The beaming factor, $F_{beam}$, is defined as $(1-\cos \theta_{jet})$, where $\theta_{jet}$ is the opening angle of the jet of the burst."871" The right hand side has à maximum value at 2,4.=12.6 with a value ≼↲↕⋅↖≺↽↔↴↓≺∢∐↓−⋟↓∪↕⋅⊡↗⊓∣⋯∶⊥⋅⊺∐∏⋟∖⇁⋅⊔∐↲≼↽⊐↥∐∏≀↧↴∐≺⇂≀↧↴↕⋅≼↲↥≀↧↴⊔∪∐↓∪↕⋅≺∢≼↲⋟∖⊽⊔∐↲∐∐⋯⋅---2.", The right hand side has a maximum value at $\textit{z}_{max} = 12.6$ with a value of $2.7 \times 10^{10} \unit{keV}^{1.43} \unit{erg}^{-1} \unit{cm}^{2}$ for $F_{beam}=1$.87213 Dp : ↽⋅ ⋅ ∙∙ ∪↓−≻⋅↙⋗⋖⊥∪↓∎∣↳↽≼↲∖ So we have a simple observational test for compliance with the Ghirlanda relation.," Thus, the Ghirlanda relation forces the limit, So we have a simple observational test for compliance with the Ghirlanda relation."873 We also have reproduced (he result that the ‘enerev ratio. for the Ghirlanda dillers from the Amati relation (e.g. Band Preece 2005)., We also have reproduced the result that the `energy ratio' for the Ghirlanda differs from the Amati relation (e.g. Band Preece 2005).874 At first 5glance. it might5 appear that this it is being5 overly A5generous to apply a beaming5 [actor Of FreeOCU=1 to ealeulate the limit of the Ghirlanda relation.," At first glance, it might appear that this it is being overly generous to apply a beaming factor of $F_{beam}=1$ to calculate the limit of the Ghirlanda relation."875 The whole point of applving such a beaming factor is to give an illustration of the lowest value the limit can have., The whole point of applying such a beaming factor is to give an illustration of the lowest value the limit can have.876 If we were apply some sort of average beaming5 factor. the value of the Ghirlanda limit would increase. resulting in many more rejected bursts.," If we were apply some sort of average beaming factor, the value of the Ghirlanda limit would increase, resulting in many more rejected bursts."877" As an example. let us use a typical jet angle. 6;,,=8.5"". which corresponds to a beaming factor. Fregn,=0.01."," As an example, let us use a typical jet angle, $\theta_{jet}=8.5^{o}$, which corresponds to a beaming factor, $F_{beam}=0.01$."878 In this case. the Ghirlancla limit would increase by a power of (wo to zz2.7xLOkeVPerg.!em?.," In this case, the Ghirlanda limit would increase by a power of two to $\approx 2.7 \times 10^{8} \unit{keV^{1.43} \; erg^{-1} \; cm^{2}}$."879" There are two reasons. however. for choosing to keep the Ghirlanda limit with a £544,=1.0."," There are two reasons, however, for choosing to keep the Ghirlanda limit with a $F_{beam}=1.0$."880 Thefirst is. the simple. mathematical. statement that we are looking. for. the maximum. value lor20 pb ποτ., Thefirst is the simple mathematical statement that we are looking for the maximum value for $\frac{E^{1.43}_{peak}}{S_{bolo}}$ .881"strong bars than SOs. still there are no galaxies having Q, z 0.5.","strong bars than S0s, still there are no galaxies having $Q_g$ $>$ 0.5."882 The conclusion in this study was that if SOs are stripped spirals. the weaker bars in SOs could indicate that bar evolution continues to proceed even after gas depletion in galaxies.," The conclusion in this study was that if S0s are stripped spirals, the weaker bars in S0s could indicate that bar evolution continues to proceed even after gas depletion in galaxies."883 Lenses appear as [lat disk components with rather sharp outer edges (43])., Lenses appear as flat disk components with rather sharp outer edges ([43]).884 However. not all lenses are directly. visible in he images.," However, not all lenses are directly visible in the images."885 In the NIRSOS Atlas (54]) lenses were generally detected as exponential subsections in the surface brightness xofiles., In the NIRS0S Atlas ([54]) lenses were generally detected as exponential subsections in the surface brightness profiles.886 NGC 524 (Fig., NGC 524 (Fig.887 6) shows all the main lens types. nuclear (nl). inner (D). and outer lens (L).," 6) shows all the main lens types, nuclear (nl), inner (l), and outer lens (L)."888 When the outer lens is very prominent. compared to the underlving disk. as in NGC 1533. it manifests as à broad bump in the surface brightness xofile. in this case having also some characteristics of a ring (tL).," When the outer lens is very prominent compared to the underlying disk, as in NGC 1533, it manifests as a broad bump in the surface brightness profile, in this case having also some characteristics of a ring (RL)."889 For clarity the examples shown of the dillerent lens tvpes are for non-barred. galaxies., For clarity the examples shown of the different lens types are for non-barred galaxies.890 However. lenses appear both in xuwred (615 and non-barred. (384) SUs. based on the classification in NIRSOS Atlas 54].," However, lenses appear both in barred $\%$ ) and non-barred $\%$ ) S0s, based on the classification in NIRS0S Atlas [54]."891 In Laurikainen et al. , In Laurikainen et al. [89252] even a arger fraction of lenses was found. but it was based on à sub-sample of NIRSOS. In barred galaxies nuclear (nl) and inner (1) lenses typically end up to the radius of the nuclear and the main bar. respectively. relating them to resonances of the rotating bar (see NGC 1543 in Fig.,"52] even a larger fraction of lenses was found, but it was based on a sub-sample of NIRS0S. In barred galaxies nuclear (nl) and inner (l) lenses typically end up to the radius of the nuclear and the main bar, respectively, relating them to resonances of the rotating bar (see NGC 1543 in Fig."893 3)., 3).894 However. not all lenses are related to resonances.," However, not all lenses are related to resonances."895 For example. series of lenses in some non-barred SOs appear. like in NGC 1411. which cannot be immediately understood in the framework of the resonance theory.," For example, series of lenses in some non-barred S0s appear, like in NGC 1411, which cannot be immediately understood in the framework of the resonance theory."896 Lenses can also be relies of significant star formation in the spiral armis as suggested in the NIRSOS Atlas (54])., Lenses can also be relics of significant star formation in the spiral arms as suggested in the NIRS0S Atlas ([54]).897 Originally the idea is [rom Bosma S]. who also confined the lens formation to the epoch of galaxy formation.," Originally the idea is from Bosma [8], who also confined the lens formation to the epoch of galaxy formation."898 Lenses are also suggested. to orm by disk instability. in à similar manner as bars 17].," Lenses are also suggested to form by disk instability, in a similar manner as bars [77]."899 Le was further suggested by Ixormendy. 43]. that bars may gradually clissolve into lenses.," It was further suggested by Kormendy [43], that bars may gradually dissolve into lenses."900 In fact there are many results in NIRSOS which are consistent with this scenario: (a) lenses in barrecl SOs often end at the bar radius. (b) S0s were found to have a smaller bar fraction. and a larger fraction of lenses than spirals 52]).," In fact there are many results in NIRS0S which are consistent with this scenario: (a) lenses in barred S0s often end at the bar radius, (b) S0s were found to have a smaller bar fraction, and a larger fraction of lenses than spirals ([52])."901 Also. (c) dissolution of bars would explain the large number of lenses in non-barred SOs in à natural manner.," Also, (c) dissolution of bars would explain the large number of lenses in non-barred S0s in a natural manner."902 Using he ellipticity of a bar. a smaller bar fraction in SOs. compared to that in spirals. was found also by Aguerri. Ménndez-AXbrevy & Corsini 4].," Using the ellipticity of a bar, a smaller bar fraction in S0s, compared to that in spirals, was found also by Aguerri, Ménndez-Abrey $\&$ Corsini [4]."903 Most. probably lenses have multiple origins. and in order to better understand their nature detailed analysis of heir dimensions ancl physical properties needs to be performed.," Most probably lenses have multiple origins, and in order to better understand their nature detailed analysis of their dimensions and physical properties needs to be performed."904 A forthcoming NURSOS paper will focus on that., A forthcoming NIRS0S paper will focus on that.905 lt is worth noticing that multiple lenses appear even in 25% of the SOs in the NIRSOS Atlas. including barred and galaxies 54]. which needs to be understood in the formation and evolution of galaxies.," It is worth noticing that multiple lenses appear even in $\%$ of the S0s in the NIRS0S Atlas, including barred and non-barred galaxies [54], which needs to be understood in the formation and evolution of galaxies."906 For example. if a large fraction of the mass in the SOs was accreted by minor mergers. it needs to be understood how the multiple lenses can survive through such processes.," For example, if a large fraction of the mass in the S0s was accreted by minor mergers, it needs to be understood how the multiple lenses can survive through such processes."907 A 2D multi-component code. BOBAR. (written by Salo. and. described in. 48]. 49D). was used. for decomposing the light distributions of the ἐν ρα images into bulges. disks. bars. ovals and lenses.," A 2D multi-component code, BDBAR (written by Salo, and described in [48], [49]), was used for decomposing the light distributions of the $K_s$ -band images into bulges, disks, bars, ovals and lenses."908 This multi-component approach turned out to be important. not only for barred galaxies. but also for galaxies with prominent lenses.," This multi-component approach turned out to be important, not only for barred galaxies, but also for galaxies with prominent lenses."909 Using artificial images this was tested by Laurikainen et al..., Using artificial images this was tested by Laurikainen et al. [910 50].,50].911 Figure Ta shows the surface brightness profile of à synthetic image with a bulge and a disk. with random noise added. whereas Th and Te show the same image after adding a small bar on top of that.," Figure 7a shows the surface brightness profile of a synthetic image with a bulge and a disk, with random noise added, whereas 7b and 7c show the same image after adding a small bar on top of that."912 Making a simple decomposition for the barred svnthetic image (in the middle) overestimates 2/7) (0.36). due to erroneous assignment of the bar Hux to the bulge. whereas the hulge-clisk-barἱ decomposition (right) recovers the correct. D/T value (01 5-02).," Making a simple bulge-disk decomposition for the barred synthetic image (in the middle) overestimates $B/T$ (=0.36), due to erroneous assignment of the bar flux to the bulge, whereas the bulge-disk-bar decomposition (right) recovers the correct $B/T$ value $B/T$ =0.27)."913 The residual image also shows a bar in the simple bulge-cisk model. but not in the bulge-clisk-bar decomposition.," The residual image also shows a bar in the simple bulge-disk model, but not in the bulge-disk-bar decomposition."914 A test for the observed NERSOS images was made by Laurikainen. Salo & Buta 49]. collected to Table 1. where 1D (using azimuthally averaged. profiles) and 2D decompositions are also compared.," A test for the observed NIRS0S images was made by Laurikainen, Salo $\&$ Buta [49], collected to Table 1, where 1D (using azimuthally averaged profiles) and 2D decompositions are also compared."915 It appears that simple bulge-clisk cdeconipositions give a similar mean D /1-ratio. independent of whether LD or 2D fitting is used. whereas the three-componoent approach gives significantly lower BY.," It appears that simple bulge-disk decompositions give a similar mean $B/T$ -ratio, independent of whether 1D or 2D fitting is used, whereas the three-component approach gives significantly lower $B/T$ ."916 The value <BYP 70.55 in the bulge-clisk decomposition is very similar to «B/Pp »=0.57 as obtained by Simien & de Vaucouleurs 67]., The value $<B/T_K>$ =0.55 in the bulge-disk decomposition is very similar to $<B/T_B>$ =0.57 as obtained by Simien $\&$ de Vaucouleurs [67].917 We also estimated that the different wavelengths used. does not. cause this difference., We also estimated that the different wavelengths used does not cause this difference.918" Adding even more components (=""final” model in the Table). like nuclear bars. further lowers the 2/7. but the change is not as dramatic as between the 2 and 3 component models."," Adding even more components (=”final” model in the Table), like nuclear bars, further lowers the $B/T$, but the change is not as dramatic as between the 2 and 3 component models."919 Phe Sérrsic index is also smaller in the bulec-cisk-bar decompositions. but even in the simplest mocels the mean value is not as large as 4. as often produced by merger simulations.," The Sérrsic index is also smaller in the bulge-disk-bar decompositions, but even in the simplest models the mean value is not as large as 4, as often produced by merger simulations."920 The 2D multi-component decompositions for the complete NIRSOS sample are discussed by Laurikainen et al..., The 2D multi-component decompositions for the complete NIRS0S sample are discussed by Laurikainen et al. [921 53]. where they are also compared. with a similar sized sample of OSUBSCS spirals. using the same decomposition approach (sce 45]. 51].,"53], where they are also compared with a similar sized sample of OSUBSGS spirals, using the same decomposition approach (see [48], [51])."922 Internal dust correction was applied to all galaxies. in a similar manner as in Graham& Worley 35].," Internal dust correction was applied to all galaxies, in a similar manner as in Graham$\&$ Worley [35]."923 For spirals these, For spirals these924The y-type and u—type distortions expected in the early Universe ( calculated in the later sections) are compared with the cosmological recombination spectrum (?) in Fig. B].,The $y$ -type and $\mu-$ type distortions expected in the early Universe ( calculated in the later sections) are compared with the cosmological recombination spectrum \citep{rcs06} in Fig. \ref{recom}.925 Clearly µ type distortions have a different spectral shape than the recombination radiation (both from hydrogen (?) and helium (?))) and y-type distortions and can be distinguished from the last two., Clearly $\mu$ type distortions have a different spectral shape than the recombination radiation (both from hydrogen \citep{Chluba2006} and helium \citep{Jose2008}) ) and $y$ -type distortions and can be distinguished from the last two.926 This is very important because the information in the µ type distortions about the early Universe physics can be extracted., This is very important because the information in the $\mu$ type distortions about the early Universe physics can be extracted.927" On the other hand, the y-type distortions from the early Universe get swamped by the much larger y-type distortions from the low redshifts and the two contributions are difficult to separate."," On the other hand, the $y$ -type distortions from the early Universe get swamped by the much larger $y$ -type distortions from the low redshifts and the two contributions are difficult to separate."928 The y distortions expected from the early Universe also have higher magnitude than the recombination spectrum in the Rayleigh-Jeans part of the spectrum but has no quasi-periodic structure like the cosmological recombination radiation., The $\mu$ distortions expected from the early Universe also have higher magnitude than the recombination spectrum in the Rayleigh-Jeans part of the spectrum but has no quasi-periodic structure like the cosmological recombination radiation.929" This flow of photons toward lower frequencies as the spectrum tries to approach the Planck spectrum due to recoil, and induced recoil is Bose-Einstein condensation of photons (??).."," This flow of photons toward lower frequencies as the spectrum tries to approach the Planck spectrum due to recoil, and induced recoil is Bose-Einstein condensation of photons \citep{is1975b,llstats}."930 We show the evolution of spectrum (solution of the coupled system of Eqs., We show the evolution of spectrum (solution of the coupled system of Eqs.931 and (16))) starting with the initial given by Eq., and ) starting with the initial given by Eq.932" with -Ysz=Yggc10? in evolution is similar to that of a spectrum with positive Ysz with the approaching a Bose-Einstein spectrum defined by n(v)=1/(e""""/sT-*""—1) with negative µ (marked in the figure).", with $-\YSZ =\YBEC=10^{-5}$ in evolution is similar to that of a spectrum with positive $\YSZ $ with the approaching a Bose-Einstein spectrum defined by $n(\nu)=1/(e^{h\nu/\kB \Te+\mu}-1)$ with negative $\mu$ (marked in the figure).933 We should emphasize that there is no singularity in the actual solutions of the Kompaneets equation plotted above., We should emphasize that there is no singularity in the actual solutions of the Kompaneets equation plotted above.934" The singularity is just in the mathematical formula, which correctly describes the spectrum at high frequencies, x>>|u|."," The singularity is just in the mathematical formula, which correctly describes the spectrum at high frequencies, $x\gg|\mu|$."935 The actual spectrum deviates from the Bose-Einstein spectrum near the singularity (positive everywhere in the Rayleigh-Jeans region) and can be described by a chemical potential decreasing in magnitude with decreasing frequency., The actual spectrum deviates from the Bose-Einstein spectrum near the singularity (positive everywhere in the Rayleigh-Jeans region) and can be described by a chemical potential decreasing in magnitude with decreasing frequency.936 The evolution at y>1 is therefore very different from the positive, The evolution at $y>1$ is therefore very different from the positive937values derived from the Schlegel extinction maps (?)..,values derived from the Schlegel extinction maps \citep{schlegel98}.938" For all of the galaxies, the upper part of the RGB is visible at Jo223 and Κοz22, and the stars above this limit are likely belonging to the IAPs of these galaxies."," For all of the galaxies, the upper part of the RGB is visible at $J_0\gtrsim23$ and $K_0\gtrsim22$, and the stars above this limit are likely belonging to the IAPs of these galaxies."939 We also overplot the 50% completeness limits in all of the panels., We also overplot the $50\%$ completeness limits in all of the panels.940" We compute the expected TRGB magnitude in both bands using the formulae given in ?,, assuming the distance moduli derived in the previous Section and themedian metallicities reported in Tab. 1.."," We compute the expected TRGB magnitude in both bands using the formulae given in \citet{valenti04}, assuming the distance moduli derived in the previous Section and themedian metallicities reported in Tab. \ref{infogen}. ."941" The resulting values are: Jo.rrgg=22.96+ 0.18, 22.560.18 and 22.560.18, and Ko,rrgg=22.00+0.18, 21.50+0.18 and 21.52+0.18 for CenA-dE1, SGC1319.1-4216 and ESO269-99, respectively."," The resulting values are: $J_{0,TRGB}=22.96\pm0.18$ , $22.56\pm0.18$ and $22.56\pm0.18$, and $K_{0,TRGB}=22.00\pm0.18$, $21.50\pm0.18$ and $21.52\pm0.18$ for CenA-dE1, SGC1319.1-4216 and ESO269-99, respectively."942" We emphasize that the TRGB is not constant as a function of metallicity in these bands, in other words its luminosity depends on the metallicity value of the galaxy and all of our targets have a considerable metallicity spread within them."," We emphasize that the TRGB is not constant as a function of metallicity in these bands, in other words its luminosity depends on the metallicity value of the galaxy and all of our targets have a considerable metallicity spread within them."943" However, as can be seen from the metallicity distribution functions presented in ?,, most of the stars in a galaxy have metallicity values around the median value, so that in the luminosity function of the galaxy the approximate locus of the TRGB will still be recognizable as a stellar count decrease toward brighter magnitudes (although not an as well-defined one, as would be the case in {-"," However, as can be seen from the metallicity distribution functions presented in \citet{crnojevic10}, most of the stars in a galaxy have metallicity values around the median value, so that in the luminosity function of the galaxy the approximate locus of the TRGB will still be recognizable as a stellar count decrease toward brighter magnitudes (although not an as well-defined one, as would be the case in $I$ -band)."944 We thus check our NIR estimates of the approximate TRGB magnitudes by additionally plotting the luminosity function for both bands in Fig. 5.., We thus check our NIR estimates of the approximate TRGB magnitudes by additionally plotting the luminosity function for both bands in Fig. \ref{lumfun}.945" The luminosity functions were dereddened, and for each magnitude bin the number of predicted Galactic foreground stars from TRILEGAL (similar to that given by Besangoon) was subtracted."," The luminosity functions were dereddened, and for each magnitude bin the number of predicted Galactic foreground stars from TRILEGAL (similar to that given by Besançoon) was subtracted."946" More precisely, we considered only a fraction of the predicted foreground stars in order to account for observational incompleteness effects (as a function of magnitude)."," More precisely, we considered only a fraction of the predicted foreground stars in order to account for observational incompleteness effects (as a function of magnitude)."947" As an example, for CenA-dEI in the magnitude bin centered at Ko=21 the completeness (as derived from artificial star tests) is ~95%, so we subtract from the dwarf galaxy star counts ~95% of the simulated TRILEGAL star counts in this magnitude bin."," As an example, for CenA-dE1 in the magnitude bin centered at $K_0=21$ the completeness (as derived from artificial star tests) is $\sim95\%$, so we subtract from the dwarf galaxy star counts $\sim95\%$ of the simulated TRILEGAL star counts in this magnitude bin."948 The expected TRGB values derived with the ? formula using the median metallicities are shown in Fig., The expected TRGB values derived with the \citet{valenti04} formula using the median metallicities are shown in Fig.949" 5 as arrows, and agree well with the observations."," \ref{lumfun} as arrows, and agree well with the observations."950" We moreover show a horizontal line including the point of origin of the arrows in order to indicate the possible range of TRGB values stemming from the range of metallicities, again computed following the ? TRGB calibration equation."," We moreover show a horizontal line including the point of origin of the arrows in order to indicate the possible range of TRGB values stemming from the range of metallicities, again computed following the \citet{valenti04} TRGB calibration equation."951" Le., if we use the lowest/highest end of a given galaxy’s metallicity range to compute the TRGB with the ? formula, we will find a correspondingly fainter/brighter TRGB value."," I.e., if we use the lowest/highest end of a given galaxy's metallicity range to compute the TRGB with the \citet{valenti04} formula, we will find a correspondingly fainter/brighter TRGB value."952" In addition, the NIR CMDs suffer from much larger incompleteness and photometric errors than the optical CMDs, such that the most metal-poor tip of the RGB is fainter than the detection limit in CenA-dEl1 and SGC 1319.1-4216, while it is close to the detection limit for ESO269-066."," In addition, the NIR CMDs suffer from much larger incompleteness and photometric errors than the optical CMDs, such that the most metal-poor tip of the RGB is fainter than the detection limit in CenA-dE1 and SGC 1319.1-4216, while it is close to the detection limit for ESO269-066."953" As a reference, in Fig."," As a reference, in Fig."954 4 we also overplot stellar isochrones over the CMDs (as in Fig. 1)), \ref{cmdnir} we also overplot stellar isochrones over the CMDs (as in Fig. \ref{cmdopt}) )955 to indicate the metallicity range of the galaxies and the shape of the TRGB., to indicate the metallicity range of the galaxies and the shape of the TRGB.956" We moreover draw a dashed line passing through the TRGB values computed from the lowest, the median and the highest metallicities found for each galaxy with the ? formula, finding a good agreement with the theoretical isochrones."," We moreover draw a dashed line passing through the TRGB values computed from the lowest, the median and the highest metallicities found for each galaxy with the \citet{valenti04} formula, finding a good agreement with the theoretical isochrones."957" Also the NIR CMDs are contaminated by Galactic foreground, but this time the luminous AGB region is not as heavily affected as it is in the optical observations (see right panel of Fig. 3))."," Also the NIR CMDs are contaminated by Galactic foreground, but this time the luminous AGB region is not as heavily affected as it is in the optical observations (see right panel of Fig. \ref{cont}) )."958 The vertical feature extending from to ~1.0 in Fig., The vertical feature extending from $J_0-K_0\sim0.3$ to $\sim1.0$ in Fig.959" 3 and 4 and over the whole magnitude range is mainly due to Galactic old disk turnoff stars (Jo—Κο 0.36), Galactic RGBand red clump stars (Jo—Ko~ 0.65), and low-mass dwarfs with Mx 0.6Mo (Jo—Ko~ 0.85, ?))."," \ref{cont} and \ref{cmdnir} and over the whole magnitude range is mainly due to Galactic old disk turnoff stars $J_0-K_0\sim0.36$ ), Galactic RGBand red clump stars $J_0-K_0\sim0.65$ ), and low-mass dwarfs with $\leq0.6$ $_{\odot}$ $J_0-K_0\sim0.85$ , \citealt{girardi05}) )."960 The fact that these stars are distributed in vertical sequences is due to therange of distances and luminosities that they span., The fact that these stars are distributed in vertical sequences is due to therange of distances and luminosities that they span.961 We, We962"channels increases from 50 at r/r4,,4,20.65 to à maximum of around 70 at r/r,,4,20.8.",channels increases from 50 at $r/r_{\rm spot}$ =0.65 to a maximum of around 70 at $r/r_{\rm spot}$ =0.8.963 In the azimuthally averaged absolute LOS velocity the observation at 30° heliocentric angle shows a large difference of 2 kms! between bg component and flow channel. whereas at 6° the sigrature of the Evershed flow is negligible.," In the azimuthally averaged absolute LOS velocity the observation at $^\circ$ heliocentric angle shows a large difference of 2 $^{-1}$ between bg component and flow channel, whereas at $^\circ$ the signature of the Evershed flow is negligible."964 In both cases the bz component shows increasing LOS velocities towards the outer boundary., In both cases the bg component shows increasing LOS velocities towards the outer boundary.965 The field strength and field inclination of both observations are similar. indicating little change of the sunspot field topology in two days.," The field strength and field inclination of both observations are similar, indicating little change of the sunspot field topology in two days."966 In both observations. the bg component is stronger by 0.5 kG in the iher and niddle penumbra. whereas at the outer boundary the strength of both components 1s nearly identical.," In both observations, the bg component is stronger by 0.5 kG in the inner and middle penumbra, whereas at the outer boundary the strength of both components is nearly identical."967 The slope with radius of the flow channel field strength decreases in the mid penumbra at. 7/rai 0.7 (dotted lines). but the field strength continues to drop with radius in both be component and flow channels.," The slope with radius of the flow channel field strength decreases in the mid penumbra at $r/r_{\rm spot}$ = 0.7 (dotted lines), but the field strength continues to drop with radius in both bg component and flow channels."968 The inclination of the flow charnel component on average exceeds 90° for t/t> 0.9. whereas the bg component never turns into horizontal fields.," The inclination of the flow channel component on average exceeds $^\circ$ for $_{\rm spot} >$ 0.9, whereas the bg component never turns into horizontal fields."969 Its maximum inclination at the outer penumbral boundary is close to 60°., Its maximum inclination at the outer penumbral boundary is close to $^\circ$.970 The temperature plot in the lower left panel shows that the radial variation of the temperature in the background component follows closely the radial curve of intensity., The temperature plot in the lower left panel shows that the radial variation of the temperature in the background component follows closely the radial curve of intensity.971 The temperature curve of the flow channel component has a similar shape with reduced amplitude. but is displaced towards the outer penumbral boundary relative to the intensity or bg temperature curve.," The temperature curve of the flow channel component has a similar shape with reduced amplitude, but is displaced towards the outer penumbral boundary relative to the intensity or bg temperature curve."972 The cecrease of the temperature in. both components at the outer penumbral boundary is presumably connected to a trade-off between stray light and temperature in the inversion., The decrease of the temperature in both components at the outer penumbral boundary is presumably connected to a trade-off between stray light and temperature in the inversion.973 At the outer boundary. the intensity level and shape of the profiles allows to use a larger stray light amount to reproduce the observed spectra. whereas in the umbra and penumbra the QS profile simply does not fit to the spectra.," At the outer boundary, the intensity level and shape of the profiles allows to use a larger stray light amount to reproduce the observed spectra, whereas in the umbra and penumbra the QS profile simply does not fit to the spectra."974 These results are 1n. good agreement with the findings of ? or ?.. both in the absolute values of. e.g.. field strength or field inclination. and in their radial variation.," These results are in good agreement with the findings of \citet{borrero+etal2004} or \citet{bellot+etal2004}, both in the absolute values of, e.g., field strength or field inclination, and in their radial variation."975 With the assumption that on large spatial scales the flow velocities in the penumbra are only due to the radially aligned Evershed flow. the LOS velocity can be decomposed into its horizontal. v4(7). and vertical component. v-(7) (e.g.222)..," With the assumption that on large spatial scales the flow velocities in the penumbra are only due to the radially aligned Evershed flow, the LOS velocity can be decomposed into its horizontal, $v_h(r)$, and vertical component, $v_z(r)$ \citep[e.g.][]{schliche+schmidt2000,bellot+balthasar+etal2003,bellot+etal2004}."976 This allows to derive theangle. a(r). re. the inclination of the flow direction to the surface. at a given radius.," This allows to derive the, $\alpha(r)$, i.e. the inclination of the flow direction to the surface, at a given radius."977 The flow angle can be derived separately for each inversion component. and then be directly compared to the average field inclination of the component. γη).," The flow angle can be derived separately for each inversion component, and then be directly compared to the average field inclination of the component, $\gamma(r)$."978 Figure 7. displays the corresponding results for the observation on 9th of August., Figure \ref{velcomp} displays the corresponding results for the observation on 9th of August.979 The flow channel component shows upflows of around | kms! (v.> 0) in the innermost penumbra. which turn into nearly horizontal flows with constant 4.6 kms! all throughout the penumbra. and finally exhibit a small downflow component for r/raj> ," The flow channel component shows upflows of around 1 $^{-1}$ $v_z > 0$ ) in the innermost penumbra, which turn into nearly horizontal flows with constant 4.6 $^{-1}$ all throughout the penumbra, and finally exhibit a small downflow component for $r/r_{\rm spot} >$ "980aud the eiissious from stars were also modeled aud removed.,and the emissions from stars were also modeled and removed.981 These attempts at loregrouud removal finally yielded values for a CIRB at 110 aad 210u.. first as a possible detection (rom an iudepeudeut group (Schevelefal.1998). and then the definitive vaue [roii the DIRBE science team (Hauserefal.1998).," These attempts at foreground removal finally yielded values for a CIRB at 140 and 240, first as a possible detection from an independent group \citep{SFD98} and then the definitive value from the DIRBE science team \citep{HAKDO98}."982. These values are shown in Fieure l iu relation to various models., These values are shown in Figure \ref{fig:dwek} in relation to various models.983 As can ye seen. the mocles are within [actors of 2-6 of the observecl values at. 140 aad 2101.. indicating a convergence between observation aud theory in he far-IR.," As can be seen, the models are within factors of 2-6 of the observed values at 140 and 240, indicating a convergence between observation and theory in the far-IR."984 Yet in the near-IR (NIB) (1-5 μι). DIRBE was unable to detect the CIRB.," Yet in the near-IR (NIR) (1 - 5 ), DIRBE was unable to detect the CIRB."985 The contamination from the zodiacal dust. as well as uucertaiutjes in the coitribution from gaaclic stars resulted in residuals wlich had large error bars and thus could ouly ye considered as upper limits 1998)., The contamination from the zodiacal dust as well as uncertainties in the contribution from galactic stars resulted in residuals which had large error bars and thus could only be considered as upper limits \citep{HAKDO98}.986. In addition. the residual maps in the NIB. [aled the tests for isotropy even in limited regious of tle sky.," In addition, the residual maps in the NIR failed the tests for isotropy even in limited regions of the sky."987 The zodiacal dust is 1iacle up of dust from asteroids an comets aud extends out to the asteroid belt., The zodiacal dust is made up of dust from asteroids and comets and extends out to the asteroid belt.988 The particles scatter aud emit light [rom the UV th‘ough the infrared with a ininimum of the sum of scattering and erission at 3.5n., The particles scatter and emit light from the UV through the infrared with a minimum of the sum of scattering and emission at 3.5.989. Iu recent «ecades a great deal has been added to our knowledge of the zodiacal dust. especialy with the advent. of satellites like IRAS and COBE.," In recent decades a great deal has been added to our knowledge of the zodiacal dust, especially with the advent of satellites like IRAS and COBE."990 Taking the sum of our knowledge today. we know that tle zodiacal dust is distributed iu a fan shape centered on the sun and that it is sliguly tilted with respect to tlie ecliHic.," Taking the sum of our knowledge today, we know that the zodiacal dust is distributed in a fan shape centered on the sun and that it is slightly tilted with respect to the ecliptic."991 There are dust bands associated with asteroidal Families. ancL there is a 2041acal dust ring at a distawe of 1.01. AU.," There are dust bands associated with asteroidal families, and there is a zodiacal dust ring at a distance of 1.01 AU."992 Unfortuiately this accumulated knowledge las not brouglt with it a compleely :iccurate model of the cdist clisribution., Unfortunately this accumulated knowledge has not brought with it a completely accurate model of the dust distribution.993 The DIRBE scieuce team bas produced a moclel of tle Cllast distribution (Ixelsalle£αἱ.1998) to subtract from the DIRBE all-sky maps. but tlie moclel parameters are uot unique axd tlis. after subtraction. the model zodiacal light leaves some resicltal ef‘ects.," The DIRBE science team has produced a model of the dust distribution \citep{KWFRA98} to subtract from the DIRBE all-sky maps, but the model parameters are not unique and thus, after subtraction, the model zodiacal light leaves some residual effects."994 The Ixelsallefad.(1995) model leaves a largee residual intensity in the eealacti€ polar caps at 25yan.. the DIRBE band that is most domiuated by the zodiacal light.," The \citet{KWFRA98} model leaves a large residual intensity in the galactic polar caps at 25, the DIRBE band that is most dominated by the zodiacal light."995 For exampe. the 25 lintensity toward the DIRBE dark spot at {Lb)(120.57.625.9*) in the DIRBE Zo«di-Subtracted Aisson Average (ZSMA) maps is 1.76srt.," For example, the 25 intensity toward the DIRBE dark spot at $(l,b) = (120.8^\circ,65.9^\circ)$ in the DIRBE Zodi-Subtracted Mission Average (ZSMA) maps is 1.76."996 This cannot be a cosmic background because ile lack of 5-ray emission toward Miku 501 limits the CIRB to be <33 ((Funuk 11998)., This cannot be a cosmic background because the lack of $\gamma$ -ray emission toward Mkn 501 limits the CIRB to be $< 33$ (Funk 1998).997 It also cannot be galactic cirrus because the 100 intensity in this field is 1.27 iiu the ZSMA maps. and Arendt ((1998) specify the ISM intensity as A(A)(1(100)—4.). with (25)=0.0180 and 44=0.66|. so the ISM intensity is 291.," It also cannot be galactic cirrus because the 100 intensity in this field is 1.27 in the ZSMA maps, and Arendt (1998) specify the ISM intensity as $R(\lambda)(I(100)-I_\circ)$, with $R(25) = 0.0480$ and $I_\circ = 0.66$, so the ISM intensity is 29."998 By elimination. most of this intensity must be zodiacal.," By elimination, most of this intensity must be zodiacal."999 Iu order to reduce the residual zodiacal enission iu the maps. Wright (1997) added one," In order to reduce the residual zodiacal emission in the maps, Wright (1997) added one"1000clusters incorporating a stellar mass function and stellar evolution.,clusters incorporating a stellar mass function and stellar evolution.1001 They concluded that extended star clusters with an initial mass of 5.8 104 Mo are sufficiently stable to survive a Hubble-time in a weak gravitational field environment., They concluded that extended star clusters with an initial mass of 5.8 $10^{4}$ $_{\sun}$ are sufficiently stable to survive a Hubble-time in a weak gravitational field environment.1002 Bekkietal.(2004) modeled the first 70 Myr of the merging of high-mass star clusters without an external tidal field., \cite{bekki04} modeled the first 70 Myr of the merging of high-mass star clusters without an external tidal field.1003 Their finding of triaxial shapes of UCDs is most likely related to the relatively short simulation time., Their finding of triaxial shapes of UCDs is most likely related to the relatively short simulation time.1004 Figure 5 illustrates that the merger objects reach a spherically symmetric shape only after a few Gyr of evolution., Figure \ref{fig_time_evol} illustrates that the merger objects reach a spherically symmetric shape only after a few Gyr of evolution.1005" Bekkietal.(2004) find a general trend of increasing reg and velocity dispersions, c, with increasing mass."," \cite{bekki04} find a general trend of increasing $r_{\rm eff}$ and velocity dispersions, $\sigma$, with increasing mass."1006 Their effective radii increase from about 8 pc at a mass of 4 106 Mo to values of the order of 20 pc at masses of 4 107 Mo., Their effective radii increase from about 8 pc at a mass of 4 $10^{6}$ $_{\sun}$ to values of the order of 20 pc at masses of 4 $10^{7}$ $_{\sun}$.1007 Figure 15 shows only those UCDs and 22419 where observations of the effective radius and the global line-of-sight velocity dispersion are available (Haseganetal.2005;," Figure \ref{sigmaandobs} shows only those UCDs and 2419 where observations of the effective radius and the global line-of-sight velocity dispersion are available \citep{hasegan,evstigneeva07,mieske08,baumgardt,hau,taylor}."1008" Figure 15aa, which shows reg versus mass, demonstrates that most of the very extended UCDs with reg> 20 pc shown in Figure 14 have so far no observed velocity dispersions."," Figure \ref{sigmaandobs}a a, which shows $r_{\rm eff}$ versus mass, demonstrates that most of the very extended UCDs with $r_{\rm eff} >$ 20 pc shown in Figure \ref{simsandobs} have so far no observed velocity dispersions."1009 Figure 15bb shows observed global line-of-sight velocity dispersions of 22419 and UCDs and of our merger objects (see Sect. 4.5)), Figure \ref{sigmaandobs}b b shows observed global line-of-sight velocity dispersions of 2419 and UCDs and of our merger objects (see Sect. \ref{sigma}) )1010 as a function of mass., as a function of mass.1011 The models have a steeper relation of σ vs. mass than the observed UCDs., The models have a steeper relation of $\sigma$ vs. mass than the observed UCDs.1012" This is due to the fact that we modeled solely extended objects with reg> 10 pc, while UCDs show a general trend of increasing reg with increasing mass (see Fig. 1))."," This is due to the fact that we modeled solely extended objects with $r_{\rm eff} >$ 10 pc, while UCDs show a general trend of increasing $r_{\rm eff}$ with increasing mass (see Fig. \ref{fig_ucdgc}) )."1013 The mean effective radii of the UCDs considered in Figure 15 increase from about 5 pcfor masses in the interval 109? to 1065 Μο to about 16 pc for masses between 10? and 107? Mo., The mean effective radii of the UCDs considered in Figure \ref{sigmaandobs} increase from about 5 pcfor masses in the interval $10^{6.0}$ to $10^{6.5}$ $_{\sun}$ to about 16 pc for masses between $10^{7.0}$ and $10^{7.5}$ $_{\sun}$ .1014" The grey dashed line in Figure 15,, which combines Eq."," The grey dashed line in Figure \ref{sigmaandobs}, which combines Eq."1015" 8 with the parameterization of re vs. mass from Dabringhausenetal.(2008),, is a good representation of the σ vs. mass relation of the observed UCDs."," \ref{Mdyn} with the parameterization of $r_{\rm eff}$ vs. mass from \cite{dabringhausen08}, is a good representation of the $\sigma$ vs. mass relation of the observed UCDs."1016 The results of Bekkietal.(2004) are closer to the observed values as their merger objects have much smaller effective radii than our objects., The results of \cite{bekki04} are closer to the observed values as their merger objects have much smaller effective radii than our objects.1017" However, it should be kept in mind that a considerable amount of UCDs with large sizes (reg> 20 pc) do not have observed velocity dispersions, yet."," However, it should be kept in mind that a considerable amount of UCDs with large sizes $r_{\rm eff} >$ 20 pc) do not have observed velocity dispersions, yet."1018 The continuous distribution of CC masses used in our parametric study results in a continuous distribution of masses of merger objects., The continuous distribution of CC masses used in our parametric study results in a continuous distribution of masses of merger objects.1019" In contrast, the observed masses of ECs and UCDs show clear accumulations near masses 10° Mo and between 107 and 10* Mo and a very low number of ECs near 10° Mo (see Fig. 14))."," In contrast, the observed masses of ECs and UCDs show clear accumulations near masses $10^5$ $_{\sun}$ and between $10^7$ and $10^8$ $_{\sun}$ and a very low number of ECs near $10^6$ $_{\sun}$ (see Fig. \ref{simsandobs}) )."1020" A straightforward interpretation on the basis of the proposed formation scenario would suggest that the mass distribution of ECs and UCDs facilitates direct conclusions on the mass spectrum of the CCs, which produced the ECs and UCDs."," A straightforward interpretation on the basis of the proposed formation scenario would suggest that the mass distribution of ECs and UCDs facilitates direct conclusions on the mass spectrum of the CCs, which produced the ECs and UCDs."1021" An interpretation. of the available data on ECs and UCDs must be done, however, with great care, as the underlying datasets (see Sect. 2))"," An interpretation of the available data on ECs and UCDs must be done, however, with great care, as the underlying datasets (see Sect. \ref{observations}) )"1022 are highly inhomogeneous and incomplete., are highly inhomogeneous and incomplete.1023" Due to the limited field of view of the Hubble Space Telescope, most extragalactic studies on GCs and ECs cover only (a part of) the optical disk of the respective galaxies."," Due to the limited field of view of the Hubble Space Telescope, most extragalactic studies on GCs and ECs cover only (a part of) the optical disk of the respective galaxies."1024" The ECs discussed in this paper are, however, halo objects located far from the optical disk of the galaxies."," The ECs discussed in this paper are, however, halo objects located far from the optical disk of the galaxies."1025" In the Milky Way, 9 out of 13 ECs have galactocentric distances greater than 20 kpc (Harris1996)."," In the Milky Way, 9 out of 13 ECs have galactocentric distances greater than 20 kpc \citep{harris}."1026". The only massive EC (Mec510° Mo and reg&20 pc) of the Milky Way, NGC2419, is located at a distance of about 92 kpc."," The only massive EC $M_{\rm EC} \approx 10^6$ $_{\sun}$ and $r_{\rm eff} \approx 20$ pc) of the Milky Way, NGC2419, is located at a distance of about 92 kpc."1027 A similar trend has been shown for the other two disk galaxies in the Local Group: 12 out of 13 ECs associated with M31 and both ECs found in M33 have projected distances well outside the optical disks of these galaxies (Huxoretal.2008;StonkutéHuxor 2009)..," A similar trend has been shown for the other two disk galaxies in the Local Group: 12 out of 13 ECs associated with M31 and both ECs found in M33 have projected distances well outside the optical disks of these galaxies \citep{huxor08,stonkute,huxor09}. ."1028" While halo ECs might be found by chance in projection to the main body of a galaxy, the probability is relatively low: if a survey coversa projected area of 20 kpc by 20 kpc and a"," While halo ECs might be found by chance in projection to the main body of a galaxy, the probability is relatively low: if a survey coversa projected area of 20 kpc by 20 kpc and a"1029the fluid interpretation.,the fluid interpretation.1030 A more complete study of perturbation growth is in progress., A more complete study of perturbation growth is in progress.1031 Of parücular interest is Cae study of the Integrated Sachs Wolle effect in (he Cosmic Microwave Dackground., Of particular interest is the study of the Integrated Sachs Wolfe effect in the Cosmic Microwave Background.1032 It is possible that the deficit of power on large angular scales (low order multipoles) max be explained in generalized Cardassian models., It is possible that the deficit of power on large angular scales (low order multipoles) may be explained in generalized Cardassian models.1033 We thank the Kavli Institute for Theoretical Physics at the University of California. santa Barbara. for hospitality.," We thank the Kavli Institute for Theoretical Physics at the University of California, Santa Barbara, for hospitality."1034 Most of this work was performed in September 2002 when all (he authors were at the Institute lor Theoretical Physics in Santa Barbara., Most of this work was performed in September 2002 when all the authors were at the Institute for Theoretical Physics in Santa Barbara.1035 Subsequent dispersal of all the authors to different parts of (he country caused a delay in the publication of (his paper., Subsequent dispersal of all the authors to different parts of the country caused a delay in the publication of this paper.1036 It is a pleasure for us to thank Arlin Crotts. Gree Tarle. the referee. ancl especially Josh Frieman lor helplul comments.," It is a pleasure for us to thank Arlin Crotts, Greg Tarle, the referee, and especially Josh Frieman for helpful comments."1037 This research was supported in part by the National Science Foundation under Grant PPIIY99-07949. NSF CAREER. grant. AST-0094335 (Y.W.). the Department of Energy grant at the University of Michigan (Ix.F. and M.LE.). the Michigan Center for Theoretical Physics (I.E. and ALL.).," This research was supported in part by the National Science Foundation under Grant PHY99-07949, NSF CAREER grant AST-0094335 (Y.W.), the Department of Energy grant at the University of Michigan (K.F. and M.L.), the Michigan Center for Theoretical Physics (K.F. and M.L.)."1038 IF. thanks the Aspen Center for Physics. where part of Chis research was conducted. for hospitality during her stay.," K.F. thanks the Aspen Center for Physics, where part of this research was conducted, for hospitality during her stay."1039The pairing heap 10)5 is a κο]-αιjusting leap that is iupleimenuted as a single heap-ordered iultiwav tree.,The pairing heap \cite{fsst} is a self-adjusting heap that is implemented as a single heap-ordered multi-way tree.1040 The basic operation on a pairius heap is the linkine operation iu which two trees are combined by liukine the root with the larger key value to the other as its leftinost child., The basic operation on a pairing heap is the linking operation in which two trees are combined by linking the root with the larger key value to the other as its leftmost child.1041 The following oerations are defined for the standard implementation of the pairiug heaps:, The following operations are defined for the standard implementation of the pairing heaps:1042"llovaiskv van der Wlis 1987) gives AV —4.75.3 mags and henee £2,,,=15.4 18.5 hes. which is consistent with that obtained by Martin (1996): 15.6 hrs.","Ilovaisky van der Klis 1987) gives $\Delta V$ =4.7–5.3 mags and hence $P_{orb}$ =15.4–18.5 hrs, which is consistent with that obtained by Martin (1996); 15.6 hrs."1043 Alasetti ct al. (, Masetti et al. (10441996) interpret a 14.7 hr periodicity as the superhump period.,1996) interpret a 14.7 hr periodicity as the superhump period.1045 Lf this interpretation is correct. then the true orbital period should be a few. percent. shorter.," If this interpretation is correct, then the true orbital period should be a few percent shorter."1046" Using equation (1) with Vo=16.3 mags and Yo 721 mags (Della Valle. Mirabel Rodriguez 1094) we estimate 2,5«18.5 hes. which is consistent with the superhump period."," Using equation (1) with $V_{O}$ =16.3 mags and $V_{Q}>$ 21 mags (Della Valle, Mirabel Rodriguez 1994) we estimate $P_{orb}<$ 18.5 hrs, which is consistent with the superhump period."1047 The amplitude of the optical outburst. A is the dillerence in magnitudes between the system in quiescence and in outburst. where Vo and Alo are the apparent anc absolute magnitudes of the system in quiescence. respectively. ane Vo and Alo are the apparent and. absolute magnitudes of the system in outburst. respectively.," The amplitude of the optical outburst, $\Delta$ is the difference in magnitudes between the system in quiescence and in outburst, where $V_{Q}$ and $M_{Q}$ are the apparent and absolute magnitudes of the system in quiescence, respectively, and $V_{O}$ and $M_{O}$ are the apparent and absolute magnitudes of the system in outburst, respectively."1048 In outburst the hot accretion disc. dominates. the optical flux me(disc)]. whereas in quiescence the observe optical flux arises from the secondary star (2) and the contribution [from the cool accretion disc.," In outburst the hot accretion disc dominates the optical flux $m_{v}(disc)$ ], whereas in quiescence the observed optical flux arises from the secondary star $m_{v}(2)$ ] and the contribution from the cool accretion disc."1049" The magnitude of the system in quiescence Vo is given by (m,(2)|2.5logf. where f is the fraction of light arising from the secondary star: f=1.0 implies all the optical [lux in quiescence come from the secondary stars."," The magnitude of the system in quiescence $V_{Q}$ is given by $m_{v}(2)+2.5\log f$, where $f$ is the fraction of light arising from the secondary star; $f$ =1.0 implies all the optical flux in quiescence come from the secondary stars."1050 Typically f. is about 0.5. Le. of the light comes from the secondary star (Chevalier llovaiskv 1989: AleClintock IHemillard 1986: and Charles 1996).," Typically $f$ is about 0.5, i.e. of the light comes from the secondary star (Chevalier Ilovaisky 1989; McClintock Remillard 1986; and Charles 1996)."1051 We can therefore write Warner (1987. 1995) finds that the luminosity of the companion stars in cataclysmic variables (L24510) are inclistinguishable from main sequence stars.," We can therefore write Warner (1987, 1995) finds that the luminosity of the companion stars in cataclysmic variables $P_{orb}$ 10) are indistinguishable from main sequence stars."1052 The secondary stars in LAINBs with orbital periods 12 hrs will also lie on the main sequence or the terminal age main sequence (Shahbaz. Navlor Charles 1997). and so we extend Warner's relation to 12 hrs.," The secondary stars in LMXBs with orbital periods 12 hrs will also lie on the main sequence or the terminal age main sequence (Shahbaz, Naylor Charles 1997), and so we extend Warner's relation to 12 hrs."1053 Hle finds that the absolute maenituce of he secondary star can be represented by Unfortunately. a similar relationship does not exist. for evolved: stars. therefore. the following is only applicable for systems with un-evolved. secondaries. Lo. svstems with ub 12 hes.," He finds that the absolute magnitude of the secondary star can be represented by Unfortunately, a similar relationship does not exist for evolved stars, therefore the following is only applicable for systems with un-evolved secondaries, i.e. systems with $P_{orb}$ 12 hrs."1054" Since we have relationships for AV and. AM,(2) as functions of the orbital period (P,s(h) 12). we can then derive a similar relationship for A(disc) As the orbital period increases. the size of the svstem. also increases."," Since we have relationships for $\Delta V$ and $M_{v}(2)$ as functions of the orbital period $P_{orb}(h)$ 12), we can then derive a similar relationship for $M_{v}(disc)$ As the orbital period increases, the size of the system also increases."1055 Lowe assume that curing outburst the size of the accretion dises in SNPs are similar ic. the accretion discs extends out to the tidal radius. then one expects the accretion disc during outburst to brighten as the orbital period. of the system increases. siniply because of. the increase in the projected surface area of the accretion disc.," If we assume that during outburst the size of the accretion discs in SXTs are similar i.e. the accretion discs extends out to the tidal radius, then one expects the accretion disc during outburst to brighten as the orbital period of the system increases, simply because of the increase in the projected surface area of the accretion disc."1056 Equation (5) can be compared with the formula give by van Paraclijs MeClintock (1904)., Equation (5) can be compared with the formula give by van Paradijs McClintock (1994).1057 They determine a relationship between the absolute magnitude of the accretion disc. X-rav luminosity and orbital period for ονΕν in outburst and LAINBs.," They determine a relationship between the absolute magnitude of the accretion disc, X-ray luminosity and orbital period for SXTs in outburst and LMXBs."1058 We can rewrite their equation as where £y is the outburst X-ray luminosity ancl Lea is the Ecldington limited luminosity for a 1.4 AL. neutron star. (, We can rewrite their equation as where $L_{X}$ is the outburst X-ray luminosity and $L_{Edd}$ is the Eddington limited luminosity for a 1.4 $_{\odot}$ neutron star. (1059Lt should. be noted that. although van Paradijs MeClintock included: the black hole candidates AQG2Z000 and GS2023|33S. removing these points does not. change he correlation significantly.),"It should be noted that, although van Paradijs McClintock included the black hole candidates A0620–00 and GS2023+338, removing these points does not change the correlation significantly.)"1060 Chen. Shrader Livio (1997) abulate logCLxfhe) for all the LAINBs and. SNTs in outburst. where in this case Levy is the Ecldineton limiting or an object with general mass AJ.," Chen, Shrader Livio (1997) tabulate $\log(L_{X}/L_{Edd})$ for all the LMXBs and SXTs in outburst, where in this case $L_{Edd}$ is the Eddington limiting for an object with general mass $M$."1061 By fitting the data points or systems with orbital periods less than 12 hrs. we obtain a linear least-squares fit of the form From equations (6) and (7) we obtain the absolute magnitude of the accretion disc as a function. of orbital period: As one can see the gradients of equation (5) and (8) are comparable (at the 90 per cent confidence level).," By fitting the data points for systems with orbital periods less than 12 hrs, we obtain a linear least-squares fit of the form From equations (6) and (7) we obtain the absolute magnitude of the accretion disc as a function of orbital period: As one can see the gradients of equation (5) and (8) are comparable (at the 90 per cent confidence level)."1062 However. unlike the correlation derived by van Paraclijs MeClintock (1994). the relation we obtain for the absolute magnitude of the accretion dise does not depend on the distance anc reddening to the SNL.," However, unlike the correlation derived by van Paradijs McClintock (1994), the relation we obtain for the absolute magnitude of the accretion disc does not depend on the distance and reddening to the SXT."1063" We therefore believe that equation (5)r is à better representation of the absolute magnitude of he aceretion disc in LMXD for svstems with £?,,,(/) 12.", We therefore believe that equation (5) is a better representation of the absolute magnitude of the accretion disc in LMXB for systems with $P_{orb}(h)$ 12.1064 Lt is interesting to compare the absolute magnitudes of he accretion discs in dwarl novae anc SAPs in outburst., It is interesting to compare the absolute magnitudes of the accretion discs in dwarf novae and SXTs in outburst.1065 Dv manipulating equation (13) of Warner (1987) for the absolute magnitude of accretion discs in dwarf novae a maximum light. we find Comparing this with equation (5). we find for a given orbital period. the accretion cdises in SNTs during outburst are more than 4 magnitudes brighter (depending on the value for f) than dwarf novae at. maximum light. which," By manipulating equation (13) of Warner (1987) for the absolute magnitude of accretion discs in dwarf novae at maximum light, we find Comparing this with equation (5), we find for a given orbital period, the accretion discs in SXTs during outburst are more than 4 magnitudes brighter (depending on the value for $f$ ) than dwarf novae at maximum light, which"1066dwarf (545-8) from optical anel X-ray moclulations at 545 and 272 seconds (see also Skillman 1996).,dwarf (545-s) from optical and X-ray modulations at 545 and 272 seconds (see also Skillman 1996).1067 The continuum spin-pulse profile consists of wo peaks and the periodogram shows most of the power on !1ο first harmonic (272-5) rathor than at the fundamental frecuceney., The continuum spin-pulse profile consists of two peaks and the periodogram shows most of the power on the first harmonic (272-s) rather than at the fundamental frequency.1068 YY Dra (Llaswell οἱ al., YY Dra (Haswell et al.1069 1997) and PQ Gem (LHellicr et al., 1997) and PQ Gem (Hellier et al.1070 1994) show similar opical clouble-pulse structure to RN JO55s|5353., 1994) show similar optical double-pulse structure to RX J0558+5353.1071 Weak X-ray double pulses are also seen in (αν Per in quiescence (Ishida ct al., Weak X-ray double pulses are also seen in GK Per in quiescence (Ishida et al.1072 1992) and in NY Avi (Ixamata anc Wovama 1993)., 1992) and in XY Ari (Kamata and Koyama 1993).1073 Phe couble-pulse structure may be explained by viewing two accreting poles through optically thin accreting regions which show their largest optical depth horizontal to the clise plane., The double-pulse structure may be explained by viewing two accreting poles through optically thin accreting regions which show their largest optical depth horizontal to the disc plane.1074 Given the insight. the emission-line. pulsations can provide in the accretion pattern. we observed. RA JO55815353 with the aim to resolve sulliciently the spin evele.," Given the insight, the emission-line pulsations can provide in the accretion pattern, we observed RX J0558+5353 with the aim to resolve sufficiently the spin cycle."1075 Preliminary results show a clouble-peak pulse in (Llarlaftis Horne 1996: Walker et al., Preliminary results show a double-peak pulse in (Harlaftis Horne 1996; Walker et al.1076 1996 also announced. pulses in the Balmer lines: Still et al., 1996 also announced pulses in the Balmer lines; Still et al.1077 1998)., 1998).1078 We observed RA J0558|5353 for 3 hours with the WIUE 4.2m at La Palma on 17 March 1995 (seeing =1 arcsecond)., We observed RX J0558+5353 for 3 hours with the WHT 4.2m at La Palma on 17 March 1995 (seeing $\approx~1$ arcsecond).1079 We used the two-arm 1815 spectrograph with Εν CCD chips to cover 63606770 with a dispersion of 19 km s+ per pixel ancl 41004900 A at 52 km , We used the two-arm ISIS spectrograph with TEK CCD chips to cover 6360–6770 with a dispersion of 19 km $^{-1}$ per pixel and 4100–4900 A at 52 km $^{-1}$ $^{-1}$.1080The 30-8 exposures were designed to sample the 545-s spin, The 30-s exposures were designed to sample the 545-s spin1081Lt is thought that the vast majority of stars form in clustered environments (withsurfacedensitiesofseveralstars.tosev-eralhundred:starspersquareparsec.e.g. 2?22)..,"It is thought that the vast majority of stars form in clustered environments \citep*[with surface densities of several stars, to several hundred stars per square parsec, e.g.][]{Lada03,Lada10,Zwart10}."1082 Whether all such clusters ave dense enough to dynamically process the primordial stellar. population is currently. the subject of debate (e.g.2)..., Whether all such clusters are dense enough to dynamically process the primordial stellar population is currently the subject of debate \citep[e.g.][]{Bressert10}.1083 However. there is observational and theoretical evidence that some clusters do at least undergo a dense phase in their evolution. a notable example being the Orion Nebula Cluster (ONC).," However, there is observational and theoretical evidence that some clusters do at least undergo a dense phase in their evolution, a notable example being the Orion Nebula Cluster (ONC)."1084 Recent work by ? has shown that the observed mass segregation in the ONC can be of a dynamical origin., Recent work by \citet{Allison09b} has shown that the observed mass segregation in the ONC can be of a dynamical origin.1085 If a ‘Luster is initially substructured (7. used. fractals to create substructure) and subvirial. then the cluster undergoes cool-collapse and the most massive stars mass segregate. in some cases forming trapezium-like svstems (7)...," If a cluster is initially substructured \citet{Allison09b} used fractals to create substructure) and subvirial, then the cluster undergoes cool-collapse and the most massive stars mass segregate, in some cases forming trapezium-like systems \citep{Allison11}."1086 Previously. it had been thought that the mass segregation in the ONC had to be primordial (2).. as the level of dynamical mass scerceation required cannot occur within  MMsyr. in clusters with smooth radial profiles.," Previously, it had been thought that the mass segregation in the ONC had to be primordial \citep{Bonnell98}, as the level of dynamical mass segregation required cannot occur within $\sim$ Myr in clusters with smooth radial profiles."1087 Given the SUCCOSS of the cool-collapse model in producing the observed levels of mass. segregation. and trapezium svstems. an investigation into the elfects of this dynamical scenario on clusters containing primordial binary populations is timely.," Given the success of the cool-collapse model in producing the observed levels of mass segregation and trapezium systems, an investigation into the effects of this dynamical scenario on clusters containing primordial binary populations is timely."1088 For simplicity. ?/— did not include primordial binaries in their simulations.," For simplicity, \citet{Allison09b} did not include primordial binaries in their simulations."1089 However. the binary fraction in the ONC is consistent with that in the field. (~cent.22)...," However, the binary fraction in the ONC is consistent with that in the field \citep[$\sim$ 45 per cent,][]{Petr98,Reipurth07}."1090 N-hodvy simulations by το... and more recentlv. 7.. have shown that in a dense cluster in virial equilibrium. a binary population with a high primordial binary fraction (~ 100 per cent) will be processed to a much lower binary fraction. consistent with the observations in the ONC.," $N$ -body simulations by \citet*{Kroupa95a,Kroupa95b,Kroupa99}, and more recently, \citet{Parker09a}, have shown that in a dense cluster in virial equilibrium, a binary population with a high primordial binary fraction $\sim$ 100 per cent) will be processed to a much lower binary fraction, consistent with the observations in the ONC."1091 ? argued that it was unlikely that the primordial binary population in the ONC was field-like. as the ONC is expanding (indicating that it was much denser in the past and. therefore had. a higher primordial binary fraction) and there are no wide (01000 aau) binary systems," \citet{Parker09a} argued that it was unlikely that the primordial binary population in the ONC was field-like, as the ONC is expanding (indicating that it was much denser in the past and therefore had a higher primordial binary fraction) and there are no wide $>$ au) binary systems"1092"where the characteristics Xi satisfies: We- apply Theorem +1.5. namely inequality: (1.7)).yo with. αἱ-and a"" replaced by WeNy and στι)Vy respectivly.. we obtain:. Usiug (6.5)). (6.13)) aud (6.15)). we compute for (fryC(0.7)«&=. and inequality (1.13)) directly follows.","where the characteristics $\o{X}_{1}^{\eps}$ satisfies: We apply Theorem \ref{theo2}, namely inequality \ref{error}) ), with $u^{\eps}$and $u^{0}$ replaced by $\o{X}^{\eps}_{1}$ and $\o{X}^{0}_{1}$ respectivly, we obtain: Using \ref{7rmhl}) ), \ref{7rmhl_9})) and \ref{7rmhl_2}) ), we compute for $(t,x)\in (0,T)\times \R^{2}$: and inequality \ref{1.12}) ) directly follows."1093 LI Application of Proposition 2.1.. and same proof as Theorci 1.5..," $\hfill{\Box}$ Application of Proposition \ref{ergo}, and same proof as Theorem \ref{theo3}."1094 The function f can be expressed as (see for instance [?])): Leta=u>0. it ∙∙is easy to check that ——————|sin2z0|di—2|sin2re|——————— ," The function $\o{f}$ can be expressed as (see for instance \cite{Picc78}) ): Let$a=-u>0$, it is easy to check that $\displaystyle1095\int^{1}_{0}\frac{dv}{a+|\sin 2\pi v|}=2\int_{|v|\leq10961/4}\frac{dv}{a+|\sin 2\pi v|}$."1097Take interested in the lamit 6. > 0andRo» o withRa > 0.," Take We are interested in the limit $a \rightarrow 0$ and $R\rightarrow1098\infty$ with $Ra\rightarrow 0$ ."1099We compute 4 n where we have: with «= =. and we have chosen Now. let e= £5. we also compute:," We compute where we have: with $c=\frac{1}{4 \pi}$ , and we have chosen Now, let $\bar{v}=\frac{v}{a}$ , we also compute:"1100Using open clusters to explore the AMR has the main advantage both in abundance and age determinations since one is dealing wilh a group of stars aud the result is less susceptible to individual errors (Carraroetal. 1998)..,Using open clusters to explore the AMR has the main advantage both in abundance and age determinations since one is dealing with a group of stars and the result is less susceptible to individual errors \citep{car98}. .1101 Cameron(1985) was the first to probe the AMR from open cluster data. and found no ageanetallicity relation based on his cluster sample.," \citet{cam85} was the first to probe the AMR from open cluster data, and found no age-metallicity relation based on his cluster sample."1102 This is not surprising since (he metallicity of the Galactic disk increased only slightly during the past 5 Gyr. while his sample of 38 clusters contained no objects older (han 5.1 Gyr.," This is not surprising since the metallicity of the Galactic disk increased only slightly during the past 5 Gyr, while his sample of 38 clusters contained no objects older than 5.1 Gyr."1103 More recently. Carraroetal.(1998). compiled a relatively homogenous sample of 37 open clusters.," More recently, \citet{car98} compiled a relatively homogenous sample of 37 open clusters."1104 The data have more expanded cluster ages wp to 9 Gvr., The data have more expanded cluster ages up to 9 Gyr.1105 After correcting for the racial abundance gradient. the derived. AMIR showed similar trend to that of nearby stars.," After correcting for the radial abundance gradient, the derived AMR showed similar trend to that of nearby stars."1106 In (his paper. we have present a new open cluster catalogue wilh much more objects.," In this paper, we have present a new open cluster catalogue with much more objects."1107 The results. based on this larger sample. would be statistically more reliable.," The results, based on this larger sample, would be statistically more reliable."1108 As we have shown in the Sect 3.2. statistically. (he space distributions (scale heights for metal poor and metal rich eroups. for voung aud old clusters) of open clusters are very likely imply (the existence of age-metallicity relation in (he Galactic disk.," As we have shown in the Sect 3.2, statistically, the space distributions (scale heights for metal poor and metal rich groups, for young and old clusters) of open clusters are very likely imply the existence of age-metallicity relation in the Galactic disk."1109 In Fig., In Fig.1110 10. we plot the dependence of metallicity on the cluster age. after correcting the radial metallicity gradient.," 10, we plot the dependence of metallicity on the cluster age, after correcting the radial metallicity gradient."1111 Unfortunately. it is difficult to draw any conclusive indication for AMIR based on this plot due to the deficiency of very older clusters.," Unfortunately, it is difficult to draw any conclusive indication for AMR based on this plot due to the deficiency of very older clusters."1112 More observational efforts should be added in findiug more older clusters., More observational efforts should be added in finding more older clusters.1113 The significant spread of the AMIR. seems real. but its origin is not vet. clear.," The significant spread of the AMR seems real, but its origin is not yet clear."1114 For the scalter in (he AMB of nearby stars. many possible catses have been suggested. such as orbital diffusion of stars. inhomogeneous chemical enrichment in the Galaxy. evolution. overlapping of different galactic substructures and so on.," For the scatter in the AMR of nearby stars, many possible causes have been suggested, such as orbital diffusion of stars, inhomogeneous chemical enrichment in the Galaxy evolution, overlapping of different galactic substructures and so on."1115 All the above mentioned effects may contribute the observed scatter. while for open clusters. the result should not be very. sensitive to orbital diffusion elfects (CorderandTwarog2001).," All the above mentioned effects may contribute the observed scatter, while for open clusters, the result should not be very sensitive to orbital diffusion effects \citep{cor01}."1116. Therefore. the scatter of AMIR along the Galactic disk from both clusters and field disk stars is an essential feature in the formation and evolution of the Milkv. Wavy.," Therefore, the scatter of AMR along the Galactic disk from both clusters and field disk stars is an essential feature in the formation and evolution of the Milky Way."1117 The main work of this paper is to compile a most complete open clusters sample with metallicity. age. distance data as well as kinematic information available.," The main work of this paper is to compile a most complete open clusters sample with metallicity, age, distance data as well as kinematic information available."1118 And upon this saniple. some statistical analvsis on spatial and metallicity distributions have been mace.," And upon this sample, some statistical analysis on spatial and metallicity distributions have been made."1119 We derived an iron radial gradient about —0.06340.008 [from the CAT 1. which is quite consistent wilh the most recent determination of oxveen gradient in nebulae and voung stars.," We derived an iron radial gradient about $-$ $\pm$ 0.008 from the CAT 1, which is quite consistent with the most recent determination of oxygen gradient in nebulae and young stars."1120 By dividing clusters into age groups. we show that iron gradient was sleeper in (he past. which is consistent with the recent result [rom Galactic planetarynebulae data.," By dividing clusters into age groups, we show that iron gradient was steeper in the past, which is consistent with the recent result from Galactic planetarynebulae data."1121population.,population.1122 However. coutiuued studies of these sub-Iuuiuous GRD-SNe iud the ISAI properties of their host ealaxies could offer important msiehts iuto the nature of these events! progenitors aud their connection with the larger sample of Επ».," However, continued studies of these sub-luminous GRB-SNe and the ISM properties of their host galaxies could offer important insights into the nature of these events' progenitors and their connection with the larger sample of LGRBs."1123 Extending these spatially-resolved ISM. studies to a wider range of nearby LORD host ealaxies would be extreuclv valuable., Extending these spatially-resolved ISM studies to a wider range of nearby LGRB host galaxies would be extremely valuable.1124 The host οslaxies of GRD 020903 (=25]. Ricker et 22002). CRB 030329/SN 2003dh ἐν=0.168. Cremer et 22003). and GRD 060218/SN 2006aj (1=0.031. Mirabal Παρα 2006) are all exceleut candidates for these studies (see also Figure 8 of Starline ct 22011).," The host galaxies of GRB 020903 $z = 0.251$, Ricker et 2002), GRB 030329/SN 2003dh $z = 0.168$, Greiner et 2003), and GRB 060218/SN 2006aj $z = 0.034$, Mirabal Halpern 2006) are all excellent candidates for these studies (see also Figure 8 of Starling et 2011)."1125 In addition. while GRB 020903 and GRB 060218 appear to belong to the same class of subhuninous GRB/SNe as CRB 100316D/SN 2010bh (a supernova association for CRB 020903. was reported im Soderbere et 22005). CRB 030320(SN 0021 has a notably higher huuinositv aud ds often cousidered to be representative of the higher-redshift “cosmological” class of LGORBs (c.e.. Stanek ct 220063: a spatially-resolved spectroscopic ISAT study. of this host galaxy could therefore prove particularly ilbluninatius.," In addition, while GRB 020903 and GRB 060218 appear to belong to the same class of subluminous GRB/SNe as GRB 100316D/SN 2010bh (a supernova association for GRB 020903 was reported in Soderberg et 2005), GRB 030329/SN 2003dh has a notably higher luminosity and is often considered to be representative of the higher-redshift “cosmological"" class of LGRBs (e.g., Stanek et 2006); a spatially-resolved spectroscopic ISM study of this host galaxy could therefore prove particularly illuminating."1126 Detailed naps of metaicity variations du all of these hosts would provide importa data on the explosion sites of LORBs. and further clarify how the metallicities of these sites colpare to ovcral host metallicities.," Detailed maps of metallicity variations in all of these hosts would provide important data on the explosion sites of LGRBs, and further clarify how the metallicities of these sites compare to overall host metallicities."1127 Additional oxoperties such as ionization parameter. star formation rato. and voung stellar x»pulation age would be useful or Characterizing the typical exosion ονπολλοιές of LOCRBDs.," Additional properties such as ionization parameter, star formation rate, and young stellar population age would be useful for characterizing the typical explosion environments of LGRBs."1128 Finally. stellar ]x»pulatiou svuthesis studies at hese sitcss ancl across f1ο host ealaxies would provide a unique opportunity to compare the differeut stellar opulatious fiat. clomiiuate the host light aud eive rise o the progenitors of these rare and intriguing events.," Finally, stellar population synthesis studies at these sites and across the host galaxies would provide a unique opportunity to compare the different stellar populations that dominate the host light and give rise to the progenitors of these rare and intriguing events."1129 We are eraο] for the hospiality and assistance of he suppyt staff at Las Cnxuias Observatorv in Chile., We are grateful for the hospitality and assistance of the support staff at Las Campanas Observatory in Chile.1130 This paper uized data from he Canuna-Ray Durst Coordinates Network (GCN) circulars., This paper utilized data from the Gamma-Ray Burst Coordinates Network (GCN) circulars.1131 EM is supported wv NASA through Einstcii Postdoctoral Fellowship eraut nunber PFO-LL0075 awarded bv the Chaudra A-av Ceouter. which is operated by the Siunithsomian Astroplivsical Observatory or NASA under contract NASS-03060.," EML is supported by NASA through Einstein Postdoctoral Fellowship grant number PF0-110075 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060."1132 GRD 1050ὲwel at Harvard is supported in part by Swift AOH eran 5OSOOTO and AOG eraut GODOGL2., GRB research at Harvard is supported in part by Swift AO5 grant 5080010 and AO6 grant 6090612.1133to the east is less closely confined around the jet than is the case on the west.,to the east is less closely confined around the jet than is the case on the west.1134 Neither the highest abundance features nor the larger near-solar region correlate with the stellar structure of NGC 6051., Neither the highest abundance features nor the larger near-solar region correlate with the stellar structure of NGC 6051.1135" The high abundance feature extends roughly across the minor axis of the galaxy, but is more extended that the eellipse."," The high abundance feature extends roughly across the minor axis of the galaxy, but is more extended that the ellipse."1136 Comparison of the maps with galaxies in the field of view shows no clear correlations., Comparison of the maps with galaxies in the field of view shows no clear correlations.1137" IC 4588, an early-type galaxy at redshift 0.051 falls at the western edge of the large cool, low abundance region to the southeast of the eastern radio lobe (region 1 in Figure 1))."," IC 4588, an early–type galaxy at redshift 0.051 falls at the western edge of the large cool, low abundance region to the southeast of the eastern radio lobe (region 1 in Figure \ref{fig:maps}) )."1138" It is possible that the cool material is associated with the galaxy, perhaps as part of a galaxy group."," It is possible that the cool material is associated with the galaxy, perhaps as part of a galaxy group."1139 ? find a small number of galaxies at approximately the same recession velocity., \citet{KoranyiGeller02} find a small number of galaxies at approximately the same recession velocity.1140" However, there is no clear surface brightness structure in the region, and there are insufficient counts to allow us to identify any additional spectral components."," However, there is no clear surface brightness structure in the region, and there are insufficient counts to allow us to identify any additional spectral components."1141" An apparent radio source coincident with IC 4588 is seen in the 610 MHz contours, but comparison with the available GMRT and VLA maps at other frequencies suggests that while there is a source at this position, its apparent extension is the result of a noise feature."," An apparent radio source coincident with IC 4588 is seen in the 610 MHz contours, but comparison with the available GMRT and VLA maps at other frequencies suggests that while there is a source at this position, its apparent extension is the result of a noise feature."1142 Figure 4 shows the map of best-fitting abundance values and the associated 90 per cent upper and lower bound maps., Figure \ref{fig:Zerr} shows the map of best-fitting abundance values and the associated 90 per cent upper and lower bound maps.1143 The central abundance feature which correlates with the jets is clear in all three maps., The central abundance feature which correlates with the jets is clear in all three maps.1144" Maps of the fit statistic show variation across the field, but do not appear correlated with the temperature or abundance maps."," Maps of the fit statistic show variation across the field, but do not appear correlated with the temperature or abundance maps."1145 This suggests that the apparent features are not the product of poor spectral fits in particular regions., This suggests that the apparent features are not the product of poor spectral fits in particular regions.1146 We test this conclusion more throughly in Section 3.2.., We test this conclusion more throughly in Section \ref{sec:bias}.1147" To examine the high abundances associated with the radio jets, we placed a number of rectangular regions along and across the jet, shown in Figure 3.."," To examine the high abundances associated with the radio jets, we placed a number of rectangular regions along and across the jet, shown in Figure \ref{fig:mapzoom}."1148" Smaller regions are used in the inner part of the jet to allow us to look for any central abundance peak, larger regions outside to minimise the uncertainties on abundance."," Smaller regions are used in the inner part of the jet to allow us to look for any central abundance peak, larger regions outside to minimise the uncertainties on abundance."1149 Spectra were extracted from these regions and fitted with an absorbed APEC model., Spectra were extracted from these regions and fitted with an absorbed APEC model.1150 The resulting abundances are shown in Figure 5.., The resulting abundances are shown in Figure \ref{fig:jetZ}.1151" The east to west profile uses the two large rectangular regions at each end, and the smaller rectangles along the jet; the north to south profile compares the upper and lower pairs of large rectangular regions, and the central region comprising the three small rectangles combined."," The east to west profile uses the two large rectangular regions at each end, and the smaller rectangles along the jet; the north to south profile compares the upper and lower pairs of large rectangular regions, and the central region comprising the three small rectangles combined."1152" While the abundances in neighbouring regions are comparable, there is a clear trend for higher abundances in the inner jets (the three central regions of the east—west profile) and declining abundance outside that area."," While the abundances in neighbouring regions are comparable, there is a clear trend for higher abundances in the inner jets (the three central regions of the east–west profile) and declining abundance outside that area."1153 The abundance of the westernmost region is lower than the abundances in the inner jet at significance., The abundance of the westernmost region is lower than the abundances in the inner jet at significance.1154" Combining regions of similar metallicity, we find that the inner part of the jets (regions 3-5 of the E-W profile, or region 3 of the N-S) is more enriched than the regions at the eastern end of the jet at 3.2c significance, but only at a 2.00 level in comparison to the western regions."," Combining regions of similar metallicity, we find that the inner part of the jets (regions 3-5 of the E–W profile, or region 3 of the N–S) is more enriched than the regions at the eastern end of the jet at $\sigma$ significance, but only at a $\sigma$ level in comparison to the western regions."1155" However, comparing the inner jet to a combination of the extreme western and eastern regions shows a 3.40 difference."," However, comparing the inner jet to a combination of the extreme western and eastern regions shows a $\sigma$ difference."1156" The northern and southern regions, combined in pairs, are less abundant at the 2.4-2.76 level, or 3c if all four are simultaneously fitted."," The northern and southern regions, combined in pairs, are less abundant at the $\sigma$ level, or $\sigma$ if all four are simultaneously fitted."1157" In general, we conclude that the high abundance region is more extended E-W than N-S, following the jet, and that its abundance is significantly greater than its surroundings, by —0.4Zo."," In general, we conclude that the high abundance region is more extended E–W than N–S, following the jet, and that its abundance is significantly greater than its surroundings, by $\sim$."1158". To test the accuracy of the maps we defined regions covering specific temperature and abundance features, extracted spectra from these regions, and fitted them."," To test the accuracy of the maps we defined regions covering specific temperature and abundance features, extracted spectra from these regions, and fitted them."1159 The regions contain between 7-660 and ~2900 net counts in the 0.7-7.0 keV band., The regions contain between $\sim$ 660 and $\sim$ 2900 net counts in the 0.7-7.0 keV band.1160" While the spectral extraction and fitting process is identical in mapping and normal spectral analysis, these regions were not constrained to contain a fixed number of counts, so should provide a test of the smoothing-like effect of the mapping process."," While the spectral extraction and fitting process is identical in mapping and normal spectral analysis, these regions were not constrained to contain a fixed number of counts, so should provide a test of the smoothing-like effect of the mapping process."1161 It also allows us to determine how well the variation within map regions corresponds to the uncertainty on the normal spectral fit., It also allows us to determine how well the variation within map regions corresponds to the uncertainty on the normal spectral fit.1162 Figure 6 shows comparisons of the range of temperatures and abundances found in the map regions with the values derived from the spectral fits., Figure \ref{fig:maptest} shows comparisons of the range of temperatures and abundances found in the map regions with the values derived from the spectral fits.1163 In general the maps appear to provide an accurate estimate, In general the maps appear to provide an accurate estimate1164the results of this work are presented.,the results of this work are presented.1165 Finally our concluding remarks are gathered in section 5., Finally our concluding remarks are gathered in section 5.1166 To obtain the spatial distribution we used archive data., To obtain the spatial distribution we used archive data.1167 We selected all-sky surveys and dedicated catalogues. which are homogeneous. have a good coverage of the galaxies. and are deep enough.," We selected all-sky surveys and dedicated catalogues, which are homogeneous, have a good coverage of the galaxies, and are deep enough."1168 We are also interested in the distribution of more specific objects. such as carbon stars. or of extended objects like star clusters and associations.," We are also interested in the distribution of more specific objects, such as carbon stars, or of extended objects like star clusters and associations."1169 The Magellanic Clouds Photometric Survey (Zaritskyetal.2002.2004) offers very rich and complete catalogues of the Magellanic Clouds. but the area they cover is somewhat limited.," The Magellanic Clouds Photometric Survey \citep{2002AJ....123..855Z,2004AJ....128.1606Z} offers very rich and complete catalogues of the Magellanic Clouds, but the area they cover is somewhat limited."1170 The MCPS for the SMC by Zaritskyetal.(2002) and for the LMC by Zaritskyetal.(2004) are catalogues of the U. B. V. and I stellar photometry of only the central 18 deg? area of the SMC and of the central 64 deg? area of the LMC.," The MCPS for the SMC by \citet{2002AJ....123..855Z} and for the LMC by \citet{2004AJ....128.1606Z} are catalogues of the U, B, V, and I stellar photometry of only the central 18 $^2$ area of the SMC and of the central 64 $^2$ area of the LMC."1171 The incompleteness becomes significant at a magnitude fainter than V < 20., The incompleteness becomes significant at a magnitude fainter than V $<$ 20.1172 This is the reason we are limiting our study to stars brighter than V=20!., This is the reason we are limiting our study to stars brighter than V=20.1173". The data are combined with 2MASS and Deep Near-Infrared Southern Sky Survey (DENIS) catalogues to provide. when available. U through K, data for the stars."," The data are combined with 2MASS and Deep Near-Infrared Southern Sky Survey (DENIS) catalogues to provide, when available, U through $_s$ data for the stars."1174 2MASS has uniformly scanned the entire sky in three near-infrared bands to detect and characterize point sources brighter than about | mJy in each band. with a signal-to-noise ratio greater than 10.," 2MASS has uniformly scanned the entire sky in three near-infrared bands to detect and characterize point sources brighter than about 1 mJy in each band, with a signal-to-noise ratio greater than 10."1175" With 2MASS we can choose as large area around the Magellanic Clouds as is necessary,", With 2MASS we can choose as large area around the Magellanic Clouds as is necessary.1176 However. the limiting magnitude is somewhat bright with J«15.8. H«15.1. and Ks«14.3 mag (Skrutskieetal.2006).," However, the limiting magnitude is somewhat bright with $<$ 15.8, $<$ 15.1, and $<$ 14.3 mag \citep{2006AJ....131.1163S}."1177 The data used here are from the 2MASS All-Sky Point Source Catalog. at IPAC Infrared Science. Archive (IRSA). Caltech/JPL (http://irsa.ipac.caltech.edu/applications/Gator).," The data used here are from the 2MASS All-Sky Point Source Catalog, at IPAC Infrared Science Archive (IRSA), Caltech/JPL (http://irsa.ipac.caltech.edu/applications/Gator)."1178 We used the criteria for the main sequence and red giant stars by Gavras(2003) (see Table 1))., We used the criteria for the main sequence and red giant stars by \citet{gavras2003} (see Table \ref{tab:age_ranges}) ).1179 These also agree with the criteria used by Nikolaev&Weinberg(2000).., These also agree with the criteria used by \citet{2000ApJ...542..804N}.1180 We used catalogues of carbon stars in the LMC and SMC from objective-prism plates. taken with the UK Schmidt Telescope.," We used catalogues of carbon stars in the LMC and SMC from objective-prism plates, taken with the UK Schmidt Telescope."1181 The LMC catalogue is produced by Kontizasetal.(2001). and the SMC catalogues are from Rebeirotetal.(1993) and Morgan&Hatzidimitriou(1995).," The LMC catalogue is produced by \citet{2001A&A...369..932K}, and the SMC catalogues are from \citet{1993A&AS...97..603R} and \citet{1995A&AS..113..539M}."1182 Hess diagrams from the various data sets are produced., Hess diagrams from the various data sets are produced.1183 A Hess diagram shows the relative density of occurrence of stars at different colour-magnitude positions in the colour-magnitude diagram., A Hess diagram shows the relative density of occurrence of stars at different colour-magnitude positions in the colour-magnitude diagram.1184 The colour-magnitude range from the data sets we use was divided in a grid with different numbers of cells for each data set., The colour-magnitude range from the data sets we use was divided in a grid with different numbers of cells for each data set.1185 The size of the cells should provide a fine grid and at the same time contain enough stars for the statistics., The size of the cells should provide a fine grid and at the same time contain enough stars for the statistics.1186 The diagrams are shown in Fig.l.., The diagrams are shown in \ref{fig:cmds}.1187 We divided the stellar content of both Magellanic Clouds into several age groups. matching features of the CMD with isochrones obtained from http://stev.oapd.inat.it/emd based on Marigoetal.(2008) and Bertellietal. (1994).," We divided the stellar content of both Magellanic Clouds into several age groups, matching features of the CMD with isochrones obtained from http://stev.oapd.inaf.it/cmd based on \citet{2008A&A...482..883M} and \citet{1994A&AS..106..275B}."1188The isochrones for 10 Myr. 100 Myr. 300 Myr. and 1 Gyr are displayed on top of the contours in the Hess diagrams.," .The isochrones for 10 Myr, 100 Myr, 300 Myr, and 1 Gyr are displayed on top of the contours in the Hess diagrams."1189 The selection criteria we use are summarized in Table 1.., The selection criteria we use are summarized in Table \ref{tab:age_ranges}.1190 To determine approximately the shape and the distribution of the galaxies we used two approaches., To determine approximately the shape and the distribution of the galaxies we used two approaches.1191 I., 1.1192 We performed star counts with a rectangular grid., We performed star counts with a rectangular grid.1193 The size of the grid cell was different for the various stellar populations., The size of the grid cell was different for the various stellar populations.1194 It was chosen in such a way as to include a sufficient number of stars and also to provide enough detail for the isodensity contour maps., It was chosen in such a way as to include a sufficient number of stars and also to provide enough detail for the isodensity contour maps.1195 The contours in these maps trace areas with equal stellar density., The contours in these maps trace areas with equal stellar density.1196 They show the overall shape of the galaxies and offer an initial insight into. their. spatial distribution., They show the overall shape of the galaxies and offer an initial insight into their spatial distribution.1197 2., 2.1198 We also used radial density profiles (RDPs) to obtain the actual distribution of the galaxies., We also used radial density profiles (RDPs) to obtain the actual distribution of the galaxies.1199 The RDPs correspond to the projected radial number density. of objects contained in concentric rings around the LMC and SMC centroids., The RDPs correspond to the projected radial number density of objects contained in concentric rings around the LMC and SMC centroids.1200 The underlying assumption for this kind of analysis is that the structures should present an important degree of radial symmetry., The underlying assumption for this kind of analysis is that the structures should present an important degree of radial symmetry.

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