We can compare Fig. 14 with the corresponding figure obtained by Mowlavi et al. (2019) in the Magellanic Clouds (see their Figs. 4 and 7 for the LMC and SMC, respectively) and in the Milky Way from DR2 data (their Fig. 10). Our conclusions here are similar: (i) There are far fewer C-rich than O-rich stars in the Galaxy than in the Clouds. The low number of carbon stars compared to O-rich stars between the Galaxy and the Clouds agrees with the decrease in TDU efficiency with increasing metallicity, and with a higher O abundance in the envelope of Galactic AGB stars on average. On the other hand, because some of the stars located in RGB and faint AGB region in Fig. 14 might be CH and/or R-hot type stars (i.e. C-rich objects, see previous sections), we could derive a lower limit for the ratio between carbon to M (O-rich) stars. It is well known that this ratio increases with the decreasing (average) metallicity of the galaxy. The primary reason for this correlation is that less C needs to be dredged-up in a metal-poor star to enable atmospheric carbon atoms to exceed those of oxygen. We derive a ratio ∼0.05, which is very similar to the average ratio derived in the disc of M 31 (see e.g. van den Bergh 2000; Hamren et al. 2015). (ii) Moreover, the distribution of O-rich stars of the Gaia-2MASS diagram covers at any given MKs magnitude a much wider range of WRP, BP − RP − WKs, J − Ks in the Galaxy than in the Clouds. The wider distribution in both axes of the O-rich zones in Fig. 14 compared with the corresponding feature in the Clouds would result from the combined effect of O-rich AGB stars turning much later into C-rich stars in the Galaxy, and the existence of a more ample range of stellar metallicities in the Galactic sample. Moreover, Fig. 14 clearly shows the high dispersion existing in MKs for the C-rich objects at a given WRP, BP − RP − WKs, J − Ks value. Part of this dispersion in MKs is compatible with the typical range in Teff (2500−3500 K) deduced for N-type stars (Bergeat et al. 2001), and to the circumstellar extinction (which we did not consider here) preferably for the objects in the extreme C-rich region. However, the mixing of carbon stars of different populations probably also contributes significantly to this dispersion in MKs (and also in Mbol). In fact, following the method outlined in Sect. 4, we have studied the kinematics of these stars and calculated the membership probability to the halo and to the thin and thick disc of the 2659 new carbon star candidates with a Vrad measurement according to EDR3 (∼40% of the sample, i.e. 1305 objects). Figure 15 shows the corresponding Toomre diagram for these carbon stars. Of this limited sample, roughly 50% belong to the thin disc (blue circles in Fig. 15), ∼30% to the thick disc or halo (blue crosses and triangles, respectively), and the rest (∼20%) have an ambiguous membership according to our membership criteria (open blue circles). Nevertheless, we note that assuming a less strict likelihood percentage to assign membership to a population (see Sect. 4), about 25% of the stars with ambiguous membership would be thick-disc and/or halo stars. Because thick-disc and halo stars are older than thin-disc stars on average, many of these C-rich objects are very probably not intrinsic AGB carbon stars, but extrinsic C-rich giants: stars with masses lower than ∼1.5 M⊙ that have become carbon rich through the mass transfer in a binary system. An alternative to this would be the possibility that the minimum mass for the formation of an intrinsic AGB stars could be as low as 1 M⊙. Some observational evidence for this can be found in the literature (Shetye et al. 2019). This conclusion is reinforced by the scale height onto the Galactic plane that can be estimated for all the C-rich stars in Fig. 14 similarly to what was done in Sect. 2: an exponential fit gives zo ∼ 600 pc, which is much larger that the scale height derived for the intrinsic N-type AGB carbon stars that clearly belong to the thin disc (see Sect. 2).
Citation: van den Bergh 2000
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