Astron. Astrophys. 317, 90-98 (1997)

5. VLM with known parallaxes

The fairly rich sample of stars with known parallaxes has represented for long time a valuable observational evidence constraining the theoretical approach to VLM structures. Fig. 8 shows the most complete CM diagram for similar objects presently available, as obtained adding recent data by Dahn et al. (1995) to the previous sample by Monet et al (1992). As it has been often suggested, the sequence of stars on the cooler edge of the observed distribution should be interpreted as the sequence of VLM stars with solar metallicity, with the additional evidence for metal poor subdwarfs spread in the hotter portion of the HR diagram. Recent results by Baraffe et al. (1995) support such a belief, disclosing that the color-magnitude location of [M/H] =-1.5 VLM structures appear in good agreement with the location of the subdwarfs in the Monet et al. (1992) sample.

As shown in the same figure, present theoretical evaluations appear in excellent agreement with such a scenario. Metal poor sequences, which represent a reasonable lower limit for the metallicity of halo subdwarfs rank indeed along the hotter boundary of the observed distribution. According to such a scenario, it has to be noticed that the CM diagram location of VLM stars appears as a metallicity indicator of unusual sensitivity, possibly to be used to investigate the metallicity distribution of stars around the Sun when the complete sample of parallaxes expected from the Hipparcos mission will be available.

Fig. 8. ( ,V-I) CM diagram for faint stars with known parallaxes from Monet et al. (1992) and Dahn et al. (1995) with superimposed theoretical distributions from the present paper for the labeled values of stellar metallicity.

As for solar metallicity VLM models, according to the results by Baraffe et al. (1995) and to the discussion given in the previous section, one expects that the treatment of the atmosphere should affect the MS location in the range 2.2 < (V-I) < 3, namely for masses in the range 0.6 < M/M 0.2. Thus, contrarily to the low metallicity case, solar metallicity models adopting relation should give a reliable picture for the extreme lower portion of the MS only. This is partially supported by data shown by Fig. 9, where we compare the recent computations presented by Baraffe et al. (1995) with our computations for solar metallicity stars, as based just on the same zero metal EOS by Saumon & Chabrier (1992) and Saumon, Chabrier & Van Horn (1995). The same figure shows the location of the similar set of computations recently presented by D'Antona & Mazzitelli (1994,1996). In all cases, theoretical sequences have been transferred into the ( ,V-I) CM diagram by adopting the results of the new grid of model atmospheres given by Allard et al.(1996) for solar metallicities only. As a whole, Fig. 9 reconfirm, as repeatedly stated in the current literature, that an accurate treatment of stellar envelopes appears a necessary ingredient to produce reliable models of the upper portion of the low main sequence of metal rich stars. However, the same figure shows that both Baraffe et al. (1995) and present computations fail in accurately reproducing the location of the lowest MS masses. As pointed out by our unknown referee, this suggests that understanding VLM structures is not complete yet, and remains a tantalizing goal.

Fig. 9. ( ,V-I) CM diagram for faint stars with known parallaxes from Monet et al. (1992) and Dahn et al. (1995) with superimposed our theoretical distributions for VLM stars of solar composition computed with the zero metal EOS of Saumon & Chabrier (1992), the theoretical models presented by Baraffe et. al. (1995) and the theoretical models of D'Antona & Mazzitelli (1994,1996).

As a final point, Fig. 9 discloses that present computations, when compared with DM, sensitively move the expected location of metal rich models, computed by integrating a relation, toward observational data. The origin of such improvement is not clear. It obviously resides either in the adopted opacity or in the EOS, if not both, with only a minor influence of the adopted transformation between theoretical and observational parameters. However, it appears difficult to understand if one out of these two physical ingredients may play a major role. Nor it is known if the error in the Magni & Mazzitelli (1979) EOS, found by Saumon (1994) and discussed in DM96 could have been of relevance for these models firstly presented in D'Antona & Mazzitelli (1994).

© European Southern Observatory (ESO) 1997

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