Value Plus Basti

[Muriel Pelley]

Thefirst paper (Pietrinferni et al. 2004; hereafter Paper I) presented scaled-solar stellar evolution models and isochrones, while Pietrinferni et al. (2006; hereafter Paper II) extended the database to α-enhanced metal compositions appropriate, for example, for modelling the stellar population of the Galactic halo (Gratton et al. 2004, and references therein). The inclusion of the asymptotic giant branch phase, and then description of the BaSTI synthetic CMD generator (SYNTHETIC_MAN) was included in Cordier et al. (2007; hereafter Paper III), whilst integrated spectra and magnitudes for the whole set of chemical compositions were published by Percival et al. (2009; hereafter Paper IV). Pietrinferni et al. (2009; hereafter Paper V) extended BaSTI to mixtures including the CNONa abundance anticorrelations observed in Galactic globular clusters (GCs), while Salaris et al. (2010; hereafter Paper VI) added to the database white dwarf (WD) cooling sequences and isochrones. As part of our ongoing effort to provide the scientific community with a self-consistent evolutionary framework for interpreting observations of a wide range of stellar populations, we now present an extension of the BaSTI archive to both extremely metal-poor and super-metal-rich stellar populations beyond the present BaSTI metallicity range.

Observations of metal-poor candidates in the HK survey (see e.g. Christlieb et al. 2008, and references therein), the Sloan Digital Sky Survey (York et al. 2000), and the SEGUE survey (Yanny et al. 2009) have vastly increased the available sample of Galactic extremely metal-poor stars, with [Fe/H]

The extension of BaSTI to both extremely metal-poor and super-metal-rich compositions will complete the creation of an updated and self-consistent theoretical evolutionary framework that covers almost the whole metallicity interval spanned by stars in both the Galaxy and extra-galactic stellar systems.

The paper is organized as follows: Sect. 2 summarizes the model physics inputs and describes calculations and main evolutionary properties. Comparisons with models available in literature are discussed in Sect. 3. A summary and final remarks follow in Sect. 4.

As for the other calculations already available in the BaSTI library, we computed models with and without convective core overshoot during the H-burning phase (when convective cores are present). We adopted exactly the same extension of the overshooting region as a function of mass that we used in our previous calculations.

The evolutionary tracks have been reduced to the same number of points, to facilitate the computation of isochrones and their use in population synthesis codes, by identifying along each evolutionary track some characteristic homologous points, also known key points (KPs) corresponding to well-defined evolutionary phases. For a careful description of the adopted KPs we refer the reader to Paper II. The whole set of evolutionary computations (except for the additional HB models discussed above) have been used to compute isochrones from 30 Myr (50 Myr for Z = 10-5) to 15 Gyr. Finally, tracks and isochrones have been transformed to various photometric systems (i.e. Johnson-Cousins, ACS, and WFC3 Vegamag) by using the same colour-Teff transformations and bolometric corrections presented in Papers I and II, and Bedin et al. (2005). These computations are made public at the BaSTI official website5.

Figure 1 reveals some important evolutionary features of these models related to their extreme metallicities. At fixed initial mass, extremely metal-poor stars are obviously much hotter and brighter than their more metal-rich counterparts. This occurrence, as is well known, is due to the huge dependence of the stellar radiative opacity on the metallicity: the larger the metallicity, the larger the radiative opacity. For an overview across the whole metallicity range spanned by the BaSTI models, Tables 1 and 2 list some relevant evolutionary features for selected models, with Z ranging from Z = 10-5 to 0.05.

From the data in Table 1 one notices that, at fixed total mass, the central H-burning lifetime is strongly affected by the metal content, monotonically increasing with increasing metallicity as a consequence of the lower brightness of the stellar structures.

Fig. 5Behaviour of MHeF with metallicity as predicted by BaSTI models for the scaled-solar metal mixture, both with (non-canonical models) and without (canonical models) MS convective core overshooting.

One important prediction of stellar model calculations is the transition mass MHeF between stars that ignite He in an electron degenerate core, and stars that enter the central He burning phase without experiencing core electron degeneracy (see e.g. Sweigart et al. 1989; 1990; Cassisi & Castellani 1993). Core mass and bolometric luminosity at He ignition change remarkably over a range of only a few tenths of solar mass, as shown in Fig. 4 for the case Z = 10-5, for models with and without MS convective core overshooting. For this reason, when discussing the evolution of stars with masses around MHeF, one often speaks of a red giant branch transition.

It is well known that the value of MHeF, for a given chemical composition, depends on the assumed efficiency of the core convective overshooting during the central H-burning stage: the larger the overshooting region, the smaller is MHeF (see Fig. 5). This is due to the larger He core mass at the end of the central H-burning stage of the models with MS core overshooting; a larger He core mass at the start of the RGB stage implies a hotter core thermal stratification, which favours the thermal conditions required for He-burning ignition. It is worth noticing how the effect of MS convective core overshooting vanishes at the lowest metallicity. This is due to the huge decrease of the size of the convective core during the central H-burning stage in very metal-poor stars, caused by the very low efficiency of the CNO cycle compared to the p-p chain.

Figure 6 shows the ZAHBs for all BaSTI scaled solar compositions. They obviously become progressively fainter and cooler with increasing metallicity, because of the previously mentioned larger envelope opacities, but mainly because of the smaller He core at He ignition. It is worth noticing the behaviour of the Z = 0.05 ZAHB, whose brightness at Teff

The generally higher central temperatures of the metal-poor models also imply that for intermediate-mass models the thermal conditions for He-burning ignition are reached earlier than for their metal-rich counterparts. Therefore, at very low Z, intermediate-mass models miss the RGB stage, as shown in Fig. 8.

Stellar models for these extreme metallicities have been calculated, so far, by few authors. Models for Z = 10-5, covering both low- and intermediate-mass stars have been computed by Cassisi et al. (1997), and can be also found as part of the Yale-Yonsei (YY) model database (see e.g. Demarque et al. 2004). Figure 4 shows the comparison between our models (without convective overshooting, as in the Cassisi et al. 1997 calculations) and Cassisi et al. (1997) results, concerning both the He core mass and surface luminosity at the RGB tip, for stellar masses around the RGB transition, while Fig. 9 compares the HR diagrams of selected H-burning models.

Figure 13 shows HR diagrams of selected HB models. There is a relatively small luminosity difference that is mainly a consequence of the different envelope He abundance, for the RGB progenitors have a similar He core mass at the RGB tip (see data in Fig. 4).

The only possibility we have to compare the Z = 0.05 models with independent calculations at exactly the same Z is to consider the models by Bressan et al. (2012, PARSEC models). The PARSEC library includes a grid point at Z = 0.05, although with a larger initial He content. The physics inputs of PARSEC calculations are similar to those adopted for the BaSTI database, the main differences being the low-T and electron conduction opacities, and some nuclear reaction rates, plus their inclusion of atomic diffusion. Their adopted scaled solar heavy element mixture starts from Grevesse & Sauval (1998), but supplemented for a subset of elements by the Caffau et al. (2011, and references therein) results. In particular, the very abundant CNO elements and Fe are amongst the metals with Caffau et al. (2011) abundances, and the resulting metal mixture is appreciably different from our calculations. Figure 15 displays a comparison of selected isochrones from BaSTI and PARSEC, from the MS to the asymptotic giant branch phase. We display the BaSTI predictions accounting for core convective overshoot, because the PARSEC calculations account for this non-canonical mixing process.

The two sets of isochrones display differences along the various branches that are not very large but are still noticeable. For isochrones populated by stars with well-developed convective cores along the MS (ages below 5 Gyr), PARSEC calculations display typically brighter and hotter TOs, while the reverse is true for the older ages. The RGB and asymptotic giant branch Teff of our models is typically larger (the reverse is true for the lower MS), and the central He burning luminosity is generally fainter. It is difficult to disentangle the various causes of these differences that are undoubtedly due to the different initial He (which however, should not appreciably affect the RGB Teff, and would in any case exacerbate the differences along this phase), the efficiency of atomic diffusion (for the older isochrone TO region only), a slightly different extension of the overshooting region (larger in the PARSEC models) in stars with well-developed MS convective cores, and the different initial metal mixture.

We have discussed the variations of several fundamental predictions of stellar evolution over the large metallicity range spanned by the full BaSTI models, and compared the new calculations with literature models at Z = 10-5 and Z = 0.05. The comparison discloses a good agreement with YY calculations from the MS to the tip of the RGB (YY models employ very similar input physics) at Z = 10-5. The existing small differences in the HR diagram are easily explained by the different initial He abundance, and the inclusion of He diffusion in YY calculations. The comparison with the Cassisi et al. (1997) Z = 10-5 models shows that the Teff difference along the RGB is due to the different low temperature radiative opacities, while it is the different initial He abundance that causes the differences in Teff and bolometric luminosity along the MS and subgiant branch. The Cassisi et al. (1997) calculations also include HB models, which tend to be slightly fainter mainly because of the lower initial Y.

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