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1 When data are available. NLSv1in our sample are radio emitting objects. though not. radio-loud. except. for LLL0323|342 which is a blazar-like NLSv1. | When data are available, NLSy1 in our sample are radio emitting objects, though not radio-loud, except for 1H 0323+342 which is a blazar-like NLSy1. |
We thank the referee. Dirk Curupe. for his useful comments and suggestions which helped. us improving our paper. | We thank the referee, Dirk Grupe, for his useful comments and suggestions which helped us improving our paper. |
We thank Alessandro Maselli for his precious contribution to the /UVOY data analysis. | We thank Alessandro Maselli for his precious contribution to the /UVOT data analysis. |
FP and RL A by XSEL/033/10/0 and ASL/INATE MNL/¢lant10/0. | FP and RL acknowledge support by ASI I/033/10/0 and ASI/INAF I/009/10/0. |
GM thanks the Spanish Ministry of Science and for support through grants AY.X2009-08059. ancl AYA2OL | GM thanks the Spanish Ministry of Science and Innovation for support through grants AYA2009-08059 and AYA2010-21490-C02-02. |
of these models). | of these models). |
Therefore, we confirm that the DLA systems could be progenitors of local dwarf galaxies. | Therefore, we confirm that the DLA systems could be progenitors of local dwarf galaxies. |
By using à spectrophotometic model coupled with a chemical evolution model, Legrand(2000) demonstrated that a continuous but very mild star formation rate (SFE as low as 107? Μοντ1) is able to reproduce the main properties of IZw 18, one of the most metal-poor BCDs we know. | By using a spectrophotometic model coupled with a chemical evolution model, \citet{Legrand00a} demonstrated that a continuous but very mild star formation rate (SFE as low as $10^{-3}$ $^{-1}$ ) is able to reproduce the main properties of IZw 18, one of the most metal-poor BCDs we know. |
Therefore, in this section we are showing model results obtained with continuous and mild SF. | Therefore, in this section we are showing model results obtained with continuous and mild SF. |
We run the same models as in Legrand(2000),, one model with a mild continuous SF (e=0.001 Gyr-1) only, and the other one with a mild continuous SF (e=0.001 ντ1) and a current burst (SFR= 0.023Moyr71, i.e., €=0.88 Gyr-! if the observed Myy=2.6x10Mo, during the last 20 Myrs). | We run the same models as in \citet{Legrand00a}, one model with a mild continuous SF $\epsilon=0.001$ $^{-1}$ ) only, and the other one with a mild continuous SF $\epsilon=0.001$ $^{-1}$ ) and a current burst $SFR=0.023$ $^{-1}$, i.e., $\epsilon=0.88$ $^{-1}$ if the observed $M_{\rm HI}=2.6\times10^7$, during the last 20 Myrs). |
The abundance ratios of different elements as predicted by these two models are shown in Fig. 18.. | The abundance ratios of different elements as predicted by these two models are shown in Fig. \ref{Fig:Legrandabund}. |
As expected, such a low continuous SF predicts a too low oxygen abundance which could not explain the majority of local BCDs. | As expected, such a low continuous SF predicts a too low oxygen abundance which could not explain the majority of local BCDs. |
In addition, by examining the evolutionary tracks in the low metallicity range, one sees that DLAs cannot be the progenitors of such galaxies. | In addition, by examining the evolutionary tracks in the low metallicity range, one sees that DLAs cannot be the progenitors of such galaxies. |
lt has long been recognised that the observed. form of the initial mass function. (IME) in the Alilky Way. imprints the stellar population with a characteristic mass somewhat below LM.. | It has long been recognised that the observed form of the initial mass function (IMF) in the Milky Way imprints the stellar population with a characteristic mass somewhat below $1 M_{\odot}$. |
The existence of a characteristic scale. linked to neither the upper nor lower limits of the observed. LME. owes itself to the observed Hattening of the mass function in this region (Scalo 1986. 1998: Ixroupa. Tout. Cilmore 1990). | The existence of a characteristic scale, linked to neither the upper nor lower limits of the observed IMF, owes itself to the observed flattening of the mass function in this region (Scalo 1986, 1998; Kroupa, Tout Gilmore 1990). |
I£ the ΙΔΗΣ is parameterised. by a series of disjoint »ower laws with respective indices à. (such that the number of stars in the mass range m to m|dm is Xm. μι). hen the total mass contributed by each power law section is dominated by the upper (lower) mass limits of the section if a is respectively < (2) 2. | If the IMF is parameterised by a series of disjoint power laws with respective indices $\alpha$ (such that the number of stars in the mass range $m$ to $m+{\rmn d}m$ is $\propto m^{-\alpha}{\rmn d}m$ ), then the total mass contributed by each power law section is dominated by the upper (lower) mass limits of the section if $\alpha$ is respectively $<$ $>$ ) $2$. |
Since a makes the transition rom 2.35 to~1.5 at around a solar mass. it follows that he characteristic stellar mass is around this value. | Since $\alpha $ makes the transition from $ \sim
2.35$ to $\sim 1.5$ at around a solar mass, it follows that the characteristic stellar mass is around this value. |
At higher masses. the ΝΗΣΟ extends in a power law for around. two orders of magnitude in mass (Salpeter 1955). | At higher masses, the IMF extends in a power law for around two orders of magnitude in mass (Salpeter 1955). |
lt is tempting to speculate on the origin of the two main features of the IME (see Larson 1995. 1996. 19982).The slope of the upper LAL power law may be set by some scale free process. | It is tempting to speculate on the origin of the two main features of the IMF (see Larson 1995, 1996, 1998a).The slope of the upper IMF power law may be set by some scale free process. |
In this regard. both coagulation and. competitive accretion. have been cited. as possible mechanisms (e.g. Alurrav Lin 1996: Bonnell et al. | In this regard, both coagulation and competitive accretion have been cited as possible mechanisms (e.g. Murray Lin 1996; Bonnell et al. |
2001). | 2001). |
The characteristic mass scale. on the other hand. is argued to be set by some ohvsical property of the star formingmo 0gas and. therefore. niv be expected to vary with environment and epoch (NTB. Larson 1998b). | The characteristic mass scale, on the other hand, is argued to be set by some physical property of the star forming gas and, therefore, may be expected to vary with environment and epoch (e.g. Larson 1998b). |
In nearby star forming regions. the characteristic stellar mass is similar to the typical Jeans mass in. dense molecular cloud. cores. | In nearby star forming regions, the characteristic stellar mass is similar to the typical Jeans mass in dense molecular cloud cores. |
This Jeans mass ds jointly se » the eas temperature and the pressure in the cores. | This Jeans mass is jointly set by the gas temperature and the pressure in the cores. |
Whereas the temperature is fixed. by the detailed: physics of heating anc cooling in molecular gas. the therma oessure in the cores is apparently in a state of rough oessure balance with the mean internal pressure of bulk motions (‘turbulence’) within the parent molecular clouc (Larson 1996. 1998a). | Whereas the temperature is fixed by the detailed physics of heating and cooling in molecular gas, the thermal pressure in the cores is apparently in a state of rough pressure balance with the mean internal pressure of bulk motions (`turbulence') within the parent molecular cloud (Larson 1996, 1998a). |
This conjecture has subsequently been supported by hyvedrodsnamic simulations of turbulent clouds. | This conjecture has subsequently been supported by hydrodynamic simulations of turbulent clouds. |
These numerical experiments show that the density. fiel eenerated by driven supersonic turbulence in an isotherma medium gives rise to à characteristic density at which the thermal pressure roughly balances the turbulent pressure in the cloud (e.g. Padoan. Nordlund Jones 1997: Padoan Nordlund 2002). | These numerical experiments show that the density field generated by driven supersonic turbulence in an isothermal medium gives rise to a characteristic density at which the thermal pressure roughly balances the turbulent pressure in the cloud (e.g. Padoan, Nordlund Jones 1997; Padoan Nordlund 2002). |
Through analysis of the clouds steady state density distribution. ancl by converting clensity to corresponding isothermal Jeans mass. these authors derive a characteristic mass ancl (approximately log-normal) EM for the eloud. | Through analysis of the cloud's steady state density distribution, and by converting density to corresponding isothermal Jeans mass, these authors derive a characteristic mass and (approximately log-normal) IMF for the cloud. |
Similar conclusions may be derived from the recent simulations of Bate. Bonnell romam (2002a.b. 2003). | Similar conclusions may be derived from the recent simulations of Bate, Bonnell Bromm (2002a,b, 2003). |
This work clillers from the above both in that the turbulence is not criven (but allowed to decay on roughly a cloud [ree-Lall time) and also. more crucially. in that it follows the collapse and fragmentation of gravitationally unstable σας down to the opacity limit ancl hence directly simulates the building up of the IME by combined fragmentation and competitive accretion. | This work differs from the above both in that the turbulence is not driven (but allowed to decay on roughly a cloud free-fall time) and also, more crucially, in that it follows the collapse and fragmentation of gravitationally unstable gas down to the opacity limit and hence directly simulates the building up of the IMF by combined fragmentation and competitive accretion. |
In this case also. the mean stellar mass appears to be set by the gas temperature and the turbulent. pressure of the initial conditions. | In this case also, the mean stellar mass appears to be set by the gas temperature and the turbulent pressure of the initial conditions. |
For observed molecular clouds with the typical | For observed molecular clouds with the typical |
The emission of a photon in the electrosphere of a quark star is the result of the scattering an electron [rom a state py to a state ptmomentumof by collision with another particle of p». | The emission of a photon in the electrosphere of a quark star is the result of the scattering of an electron from a state $\vec{p}_{1}$ to a state }$ by collision with another particle of momentum $\vec{p}_{2} $. |
In order to calculate the photon emissivity due to electron collisions in the electrosphere of quark stus we consider (he rate of collisions ΝΔ in which a soft photon of energy w is emitted. | In order to calculate the photon emissivity due to electron collisions in the electrosphere of quark stars we consider the rate of collisions $dN$ in which a soft photon of energy $\omega $ is emitted. |
According to the general principles of quantum electrodvnamies therate of collisions is given by (Berestetskiietal.1932) where is (he scattering rate. | According to the general principles of quantum electrodynamics therate of collisions is given by \citep{La}
where is the scattering rate. |
We have denoted bv (e;.p;) the i-th electron energy and momentum. where e;=(p?4-m7)5.1/2 . and by (s.i) the photon energy and momentum. f([e(p).— is th | We have denoted by $\left( \epsilon _{i},\vec{p}%
_{i}\right) the i-th electron energy and momentum, where $\epsilon
_{i}=\left( p_{i}^{2}+m^{2}\right) ^{1/2}$ , and by $\left( \omega ,\vec{k}%
\right) the photon energy and momentum. $f(\left[ \epsilon (p)-\mu _{e}%
\right] =1/\left( \exp \left[ \frac{\epsilon (p)-\mu
_{e}}{T}\right]
+1\right) |
e Fermi-Dirac distribution function for the i-th electron. and M is Che scattering amplitude. | is the Fermi-Dirac distribution function for the i-th electron \citep{Ch68} and $M$ is the scattering amplitude. |
Ii order to obtain an explicit expression for the scaltering rate we adopt the method developed in Berestetskiietal.(1982).. bv. assuming that the photon emission process is quasi-classical. | In order to obtain an explicit expression for the scattering rate we adopt the method developed in \citet{La}, by assuming that the photon emission process is quasi-classical. |
Consequently. for soft photon emission α can be written in a factorized form. dM=αμαΕς with dM being the elastic scattering rate and ΤΕ. the probability of emission of a photon of energy w by any dMcanofthescatteredelectrons. | Consequently, for soft photon emission $dW$ can be written in a factorized form, $dW=dW_{0}dW_{\gamma }$, with $dW_{0}$ being the elastic electron-electron scattering rate and $dW_{\gamma }$ the probability of emission of a photon of energy $\omega $ by any of the scattered electrons. |
beobtained from where AM, is the electron-electron scattering amplitude. | $%
dW_{0 can beobtained from where $M_{0}$ is the electron-electron scattering amplitude. |
The reaction rate can be expressed in a more convenient form in terms of the cross section σ of the reaction. | The reaction rate can be expressed in a more convenient form in terms of the cross section $\sigma $ of the reaction. |
In terms of the cross section σ of the emission process (he rate of collisions in the electron plasma in which a photon of energy v<<e is emitted can be obtained as (Chin1963) where dO(0.6) isthe solid angle element in the direction (8. 6). αἱ.go are the statistical weights and (Pj. 0$ are the velocities of the interactingparticles. | In terms of the cross section $\sigma $ of the emission process the rate of collisions in the electron plasma in which a photon of energy $%
\omega <<\epsilon is emitted can be obtained as \citep{Ch68}
where $d\Omega \left( \theta ,\phi \right) $ isthe solid angle element in the direction $\left( \theta ,\phi \right) $ , $g_{1},g_{2}$ are the statistical weights and $\vec{v}_{1}$ , $\vec{v}_{2}$ are the velocities of the interactingparticles. |
The invariant relative velocity | The invariant relative velocity |
sections lor jo=5 and 20 [or collisions with para-Ils. | sections for $j_2=5$ and 20 for collisions with $_2$. |
A number of resonances are evident between 7 and LO? + which result in significant modulation in the rate coellicients. | A number of resonances are evident between $^{-2}$ and $^{2}$ $^{-1}$ which result in significant modulation in the rate coefficients. |
llowever. (he magnitude of (he resonances are seen to decrease willi jo. | However, the magnitude of the resonances are seen to decrease with $j_2$ . |
Cross sections for some other jo can be found in Yangetal.(2006a.b).. | Cross sections for some other $j_2$ can be found in \citet{yang061,yang062}. |
The agreement between the CC and CS caleulations is shown to be excellent with a difference (vpically better than ~20%.. juslilving the adoption of the CS approximation at the higher energies. | The agreement between the CC and CS calculations is shown to be excellent with a difference typically better than $\sim$, justifying the adoption of the CS approximation at the higher energies. |
Quenching rate coellicients from initial rotational states. jo=5. 10. 20. and 40. al temperatures ranging Irom 1 to 3000 Ix are shown in Fies. | Quenching rate coefficients from initial rotational states, $j_2$ =5, 10, 20, and 40, at temperatures ranging from 1 to 3000 K are shown in Figs. |
2-0forC'Oscalteringwilhparea andorltho H3. | \ref{fig2}$ $-$ \ref{fig9} for CO scattering with para- and $_2$. |
Unfortunately. we are unaware of anv experimental rate coefficient data for (hese initial states. | Unfortunately, we are unaware of any experimental rate coefficient data for these initial states. |
Therefore. (he current resulis are compared with the theoretical rate coellicients of Flower(2001).. which were obtained over the limited temperature range. 5 {ο 400 Ix. with the Vox PES. | Therefore, the current results are compared with the theoretical rate coefficients of \citet{flower01}, which were obtained over the limited temperature range, 5 to 400 K, with the $_{98}$ PES. |
Similar results with Vy; were obtained by Wernlietal.(2006) for rotational levels of CO up to 5 and temperatures in (he range 5—70 Ix. Wernli also presented an analvlic Π valid in the same temperature range. | Similar results with $_{04}$ were obtained by \citet{wern06} for rotational levels of CO up to 5 and temperatures in the range $-$ 70 K. Wernli also presented an analytic fit valid in the same temperature range. |
The rate coefficients caleulated using their analytic relation are almost identical to the current results. | The rate coefficients calculated using their analytic relation are almost identical to the current results. |
For the quenching of jo=5 due to para-Ll» collisions. Fig. | For the quenching of $j_2$ =5 due to $_2$ collisions, Fig. |
2. shows that for temperatures between ~1 and LOO I. which is the van der Waals interaction-dominated regime. the rale coefficients exhibit an undulatory temperature dependence for Ajo=—1. —2. and —3 transitions due to the presence of though the magnitude of the undulations decrease with increasing jo. | \ref{fig2} shows that for temperatures between $\sim$ 1 and 100 K, which is the van der Waals interaction-dominated regime, the rate coefficients exhibit an undulatory temperature dependence for $\Delta j_2=-1$, $-2$, and $-3$ transitions due to the presence of though the magnitude of the undulations decrease with increasing $j_2$. |
Al temperatures above ~100 Ik. the rate coellicients generally increase wilh increasing temperature. | At temperatures above $\sim$ 100 K, the rate coefficients generally increase with increasing temperature. |
Comparison wilh (he rate coefficients of Flower shows that al temperatures higher than ~50 Ix. there is generally good agreement. except Lor Js=0 and 2. where Flowers resulis get larger αἱ temperatures above about 200 Ix. The | Comparison with the rate coefficients of Flower shows that at temperatures higher than$\sim$ 50 K, there is generally good agreement, except for $j_2'=0$ and 2, where Flower's results get larger at temperatures above about 200 K. The |
tests for distinguishing between dillerent cosmogonies (see. e.g... White. Efstathiou Frenk 1993: Eke. Cole Frenk 1996: DBartelmann et al. | tests for distinguishing between different cosmogonies (see, e.g., White, Efstathiou Frenk 1993; Eke, Cole Frenk 1996; Bartelmann et al. |
1998: Boreani οἱ 11998. and references therein). their mass-basecl detection would indeed be of great interest. | 1998; Borgani et 1998, and references therein), their mass-based detection would indeed be of great interest. |
masseselected sample of haloes may lead to the detection of clusters with very faint emission which could be missed by other selection criteria. | A mass-selected sample of haloes may lead to the detection of clusters with very faint emission which could be missed by other selection criteria. |
The basic method cliscussed in 896 Is to use the aperture mass AZ44(0) technique (Ixaiser. 1995: Squires Ixaiser 1996) on deep wide-field images. | The basic method discussed in S96 is to use the aperture mass $M_{\rm
ap}(\theta)$ technique (Kaiser 1995; Squires Kaiser 1996) on deep wide-field images. |
The aperture mass is the projected density field of the mass inhomogencitics between us and the population of faint high-redshift galaxies. weighted by a redshift-dependent term and filtered through a function of zero net weight (c... a Mexican. hat). | The aperture mass is the projected density field of the mass inhomogeneities between us and the population of faint high-redshift galaxies, weighted by a redshift-dependent term and filtered through a function of zero net weight (e.g., a Mexican hat). |
The advantage of this measure is that it can be expressed directly in terms of the shear. for which the observed image ellipticities provide an unbiased estimate. | The advantage of this measure is that it can be expressed directly in terms of the shear, for which the observed image ellipticities provide an unbiased estimate. |
Thus. an estimate or the aperture mass can be expressed directly in terms of observables. with well defined signal-to-noise ratio. | Thus, an estimate for the aperture mass can be expressed directly in terms of observables, with well defined signal-to-noise ratio. |
Hence. a (dark) matter concentration would be ‘seen’ as a high S/N »vak in the aperture mass map. | Hence, a (dark) matter concentration would be `seen' as a high S/N peak in the aperture mass map. |
In this paper. we investigate the statistics. of such »eaks in various cosmological models. | In this paper, we investigate the statistics of such peaks in various cosmological models. |
The number density of haloes is calculated. using the Press-Schechter. (1974) ormalism. and their density. profile is approximated by the universal halo profile found by Navarro. Erenk White (1996. 1997: hereafter NEW). | The number density of haloes is calculated using the Press-Schechter (1974) formalism, and their density profile is approximated by the universal halo profile found by Navarro, Frenk White (1996, 1997; hereafter NFW). |
In. 22 we summarize our method. and estimate signal-to-noise statistics in 33. | In 2 we summarize our method, and estimate signal-to-noise statistics in 3. |
Phe number of haloes of given Mj(9). as à function of filter scale 8. and source and lens redshift. is derived. in 44. | The number of haloes of given $M_{\rm ap}(\theta)$, as a function of filter scale $\theta$, and source and lens redshift, is derived in 4. |
We discuss the degree to which observations can be used to distinguish between these various cosmologies in 55. ancl present our conclusions in 66. | We discuss the degree to which observations can be used to distinguish between these various cosmologies in 5, and present our conclusions in 6. |
Following SOG. we define the spatially filbered mass inside a circular aperture of angular radius 6. where the continuous weight function C(9) vanishes for 0c8. | Following S96, we define the spatially filtered mass inside a circular aperture of angular radius $\theta$, where the continuous weight function $U(\vartheta)$ vanishes for $\vartheta>\theta$. |
IW) is a compensated filter function. one can express Ay in terms of the tangential shear inside the circle where is the tangential component of the shear at. position Y=(0cos@,9sin$). and the function Q is related to U hv ( | If $U(\vartheta)$ is a compensated filter function, one can express $M_{\rm ap}$ in terms of the tangential shear inside the circle where is the tangential component of the shear at position $\mbox{\boldmath$ $}=(\vartheta \ \mbox{cos} \ \phi,\vartheta \ \mbox{sin} \ \phi)$, and the function $Q$ is related to $U$ by We use a filter function from the familiy given in Schneider et al. ( |
5336 use | 1998), specifically we choosethe one with $l=1$. |
a f | Then writing $U(\vartheta)=u(\vartheta/
\theta)/\theta^2$, and $Q(\vartheta)=
q(\vartheta/\theta)/\theta^2$, and with $u(x)=0$ and $q(x)=0$ for $x>1$. |
ilter fur and @=r.δη. | We willdescribe the mass density of dark matter haloes with the universal density profile introduced by NFW, with $\Omega_{\rm d}$ and $\Omega_{\rm v}$ denote the present day density parameters in dust and in vacuum energy respectively. |
Va Dy is the angular\ | Haloes identified at redshift $z$ with mass $M$ are described by the characteristic density contrast $\delta_{\rm c}$ and the scaling radius $r_{\rm s}=r_{200}/c$ where $c$ is the concentration parameter (which is a function of $\delta_{\rm c}$ ), and $r_{200}$ is the virial radius defined such that a sphere with radius $r_{200}$ of mean interior density $200 \ \rho_{\rm crit}$ contains the halo mass $M_{200}$. |
dia | We compute the parameters which specify the NFW profile according to the description in NFW using the fitting formulae given there. |
meter d | The surface mass density of the NFW-profile is given by (see Bartelmann 1996) with and $\theta_{\rm s}= r_{\rm s}/D_{\rm d}$. |
istance to the lens. | $D_{\rm d}$ is the angular diameter distance to the lens. |
Introducing the critical surface density with 2, and 24, being the angular diameter distances to the source and from the lens to the source. we define the dimensionless surface mass density (convergence) which is a function of source recshift with The second important quantity for lensing ellects is the complex shear defined by where vc is given by the two-dimensional Poisson equation In the case of an axi-svmumetrie density profile. the magnitude of the shear is given by where We obtain with fore2 l. and [orar« 1. | Introducing the critical surface density with $D_{\rm s}$ and $D_{\rm ds}$ being the angular diameter distances to the source and from the lens to the source, we define the dimensionless surface mass density (convergence) which is a function of source redshift with The second important quantity for lensing effects is the complex shear defined by where $\psi$ is given by the two-dimensional Poisson equation In the case of an axi-symmetric density profile, the magnitude of the shear is given by where We obtain with $\mbox{for} \ x>1$ , and $\mbox{for} \ x<1$ . |
Xccording to eq.(4) the tangential shear is We assume a normalized source redshift clistribution of the form | According to eq.(4) the tangential shear is We assume a normalized source redshift distribution of the form |
al. 2010.. | al. \cite{dotter}, |
for a recent analysis and discussion). | for a recent analysis and discussion). |
It should be noted that the age estimates provided by POS are consistent within ] c uncertainties with this conclusion. | It should be noted that the age estimates provided by P05 are consistent within 1 $\sigma$ uncertainties with this conclusion. |
However. these clusters were singled out as likely younger than the bulk of M31 clusters at similar metallicity and this classification. is clearly not confirmed here. | However, these clusters were singled out as likely younger than the bulk of M31 clusters at similar metallicity and this classification is clearly not confirmed here. |
On the other hand. ages as young as claimed by W110 and FIO (all €4 Gyr) are clearly ruled out by the observed CMDs for all main targets of the present study except (possibly) BOS8. | On the other hand, ages as young as claimed by W10 and F10 (all $\le 4$ Gyr) are clearly ruled out by the observed CMDs for all main targets of the present study except (possibly) B058. |
The comparison with the results by BeOS significantly depends on the considered set of age estimates provided by these authors (see Table 2)). | The comparison with the results by Be05 significantly depends on the considered set of age estimates provided by these authors (see Table \ref{tab:age}) ). |
It may be interesting to note that except for B350. their estimates based on the BCO3 SPSS models are not compatible (too young) with the observed CMDs. even taking into account the reported uncertainties; on the other hand. the estimates obtainec from the TO3 models are similar to those by POS. and. in particular. the old age reported for B350 fully agrees with our results. | It may be interesting to note that except for B350, their estimates based on the BC03 SPSS models are not compatible (too young) with the observed CMDs, even taking into account the reported uncertainties; on the other hand, the estimates obtained from the T03 models are similar to those by P05, and, in particular, the old age reported for B350 fully agrees with our results. |
The sensitivity of colors and/or spectral indices generally used as age diagnosties (like. for example. Hf) to the morphology of the horizontal branch is known for long time: for example. it was already taken into account in the models of the integrated light of SSPs by Buzzoni (1989)). | The sensitivity of colors and/or spectral indices generally used as age diagnostics (like, for example, $\beta$ ) to the morphology of the horizontal branch is known for long time: for example, it was already taken into account in the models of the integrated light of SSPs by Buzzoni \cite{buz}) ). |
Bright and hot BHB stars typical of an old population may mimic the effect of a hotter MSTO. r.e.. of a younger age (see Lee et al. | Bright and hot BHB stars typical of an old population may mimic the effect of a hotter MSTO, i.e., of a younger age (see Lee et al. |
2000 for a thorough discussion. and Pereival Salaris 2010. for a very recent investigation in the context of SPSS modeling). | \cite{lee} for a thorough discussion, and Percival Salaris \cite{perci} for a very recent investigation in the context of SPSS modeling). |
It is tempting to identify this as the main reason at the origin of the erroneous classification of B292 and B350 as intermediate-age clusters (Rey et al. 2007 )). | It is tempting to identify this as the main reason at the origin of the erroneous classification of B292 and B350 as intermediate-age clusters (Rey et al. \cite{rey07}) ). |
To gain a deeper insight into this hypothesis. we considered all M31 GCs for which a safe classification of their HB morphology can be obtained from their CMDs. taken from different sources (see Sect. | To gain a deeper insight into this hypothesis, we considered all M31 GCs for which a safe classification of their HB morphology can be obtained from their CMDs, taken from different sources (see Sect. |
| for an exhaustive list). | \ref{int} for an exhaustive list). |
We simply divided the clusters into BHB (blue squares) and RHB (red circles) from the inspection of their CMD: our classification criterion is illustrated in Fig. 11.. | We simply divided the clusters into BHB (blue squares) and RHB (red circles) from the inspection of their CMD; our classification criterion is illustrated in Fig. \ref{cmd6}. |
We ended up with a sample of 33 clusters that also have estimates of the HB and Mg? indices reported in the RBC (see GO9). | We ended up with a sample of 33 clusters that also have estimates of the $\beta$ and Mg2 indices reported in the RBC (see G09). |
The classical age-diagnostic plot Mg? vs. Hf for these clusters is shown in the upper panel of Fig. 12: | The classical age-diagnostic plot Mg2 vs. $\beta$ for these clusters is shown in the upper panel of Fig. \ref{HbMg2}; ; |
a grid of SSP models from the most recent set by Thomas et al. (2010)) | a grid of SSP models from the most recent set by Thomas et al. \cite{thomas}) ) |
is over-plotted as a reference. | is over-plotted as a reference. |
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