For the supergiants, this effect becomes stronger in the 40 to 45 kK regime, while for S50, finally, the impact of X-rays on the Civlines becomes weak again, presumably because in this temperature range the stellar radiation field dominates in controlling the ionization equilibrium. We note that for the D50 dwarf model the changes remain considerable though.
In Fig. 3.15, we detail this behavior, for our S40 model, where the effect is strongest. On the figure we plot a model without X-rays, a model with intermediate shock radiation, and the model with our strongest X-ray emission. The dotted profiles give an impression of a corresponding decrease in carbon abundance which would be necessary to mimic the X-ray effect, which is 0.3 and 0.6 dex, respectively. The other way round, for stars that have been analyzed without X-rays but exhibit a strong X-ray radiation field, the originally derived carbon abundance might need to be increased by such an amount to compensate for the missing ray field. Our investigation clearly indicates that X-rays may be important for the Civanalysis of supergiant stars with temperatures around 40 to 45 kK (e.g., the prototypicalζPup) and for (very) hot dwarfs, in particular if no lines from other carbon ions are present.
In summary, the changes are marginal for not too hot dwarfs, and affect only a few Ciiilines (the triplet at 4647-50-51, the doublet at 4663-65, and Ciii4186, 5272, 5696, 5826, 6744, 8500 ˚A) which might be used with a lower weight in abundance analysis.
In contrast, and at least for supergiants in the range between 40 to 45 kK, various lines become substantially modified when accounting for strong emission from wind-embedded shocks, in particu-lar the two Ciiilines that are strongly coupled to the UV, and the Civlines. The potential differences in abundances derived from these lines (∼0.1 dex from Ciii4647-50-51 and Ciii5696, and ∼0.3 to 0.6 dex from Civ5801-5811) may complicate the analysis considerably, and we conclude that the carbon analysis of supergiants earlier than O7 should include X-ray radiation using typical default values, as already standard for CMFGEN modeling. We note (i) that this problem might have also affected our analysis of HD 169582, and (ii) that X-rays might need to be considered in the analysis of (very) hot dwarfs as well, due to their impact on Civ.
de-Figure 3.15: Impact of X-ray radiation on the Civ5801-12 lines, for a supergiant model withTeff
= 40 kK (S40). Black: Line profiles for the “standard” model without shocks; green: model with an intermediate X-ray strength; red: model with strong X-ray radiation. The dotted profiles refer to alternative, no-shock models with reduced carbon abundance, mimicking the effect of shock radiation.
We note that the displayed sensitivity is strongest for the considered lines and parameter range, see text.
scription of Ciiand Ciiiin the wind, for almost all temperatures17. When DR is neglected, we find less recombination from Civto Ciii, and consequently also less Cii. Typical differences amount to 1 dex. Similar results were reproduced with WM-basic, though we found an important difference between the underlying data: The strengths of the stabilizing transitions used in WM-basic are typ-ically larger (up to a factor of approximately two) than those corresponding to the strengths of the resonances included in the OPACITY Project data we have used, leading to stronger recombination in WM-basicmodels. Even though this potential uncertainty will not affect the majority of optical lines formed in the photosphere, it can be influential in upcoming analysis of UV wind lines. In the end, we prefered to use the OPACITY Project data, for reasons outlined in Sect. 3.3.1.
Unfortunately, optical carbon lines are rare, and often blended by other lines, which hampers
17except forTeff>∼50 kK
the spectral diagnostics, particularly for largevsiniand/or vmac. Thus, we compiled and selected a maximum set of potential diagnostic lines, including blended ones and also those with a complex formation mechanism controlled by UV transitions (as decribed by Martins et al. 2012).
Since the majority of metal lines are weak, they are very sensitive to comparatively small varia-tions of stellar parameters, and this is particularly true for carbon lines. A change of±1000 to 1500 K in effective temperature, or±0.2 dex in logg, results in considerable changes of line strength. Addi-tionally, some of the lines are also sensitive to ˙Mvariations. Mainly in supergiants (and hot dwarfs), due to their dense winds, a decrease in ˙Mby a factor of three produces an effect stronger than a de-crease of 1500 K inTeff or an increase of 0.2 dex in logg. Thus one needs to precisely determine ˙M by reproducing, for example Hαand Heii4686, before aiming at a carbon analysis.
As a first application of our “new” capability to analyze carbon lines by means of FASTWIND, using our newly developed model atom, we investigated the spectra of a sample of six O-type dwarfs and supergiants, kindly provided to us by Holgado et al. (2018). We first convinced ourselves that the stellar and wind parameters derived by Holgado et al. (2018) are reproduced by our own H and He analysis. Thereafter, we varied Teff and logg inside a range where the fit quality of the H/He lines could be preserved (roughly,±1000 K and±0.1 dex, respectively), while fitting the carbon lines from different ions in an optimum way, with an (almost) unique abundance in most cases. The finally quoted abundance was taken, if necessary, from the best-compromise solution, and from the reaction of the carbon lines when varying the abundances, we obtained a rough estimate on the corresponding error, typically in the range±0.2 to±0.3 dex.
Only for the coolest (O9.7) dwarf, lines from all three ionization stages are present in parallel, and our analysis resulted in a satisfactory reproduction of the ionization balance. In most other stars, only Ciiiand Civare visible in parallel. For one object withTeff≈40 kK (HD 169582), these Ciii/iv lines disagree by∼0.2 dex in [C/H], which might be related to the influence from X-ray radiation (see below). The hottest (O3) supergiant, CygOB2-7, displays only Civ, which renders the analysis quite uncertain (and there are only two suitable Civlines in the optical).
We were, nevertheless, able to derive carbon abundances for all the considered objects, and found in most cases a moderate depletion compared to the solar value by Asplund et al. (2009), except for CygOB2-7, where a larger depletion by 0.4 dex was derived (though more uncertain than the other values). Two of our cool objects had already been analyzed by Martins et al. (2015a). For both of them we confirmed rather similar stellar parameters, but only for one of them (a dwarf) also the carbon abundances are quite similar, while for the other (a supergiant), there is an only marginal overlap within the errors.
To obtain further constraints on the reliability of our new set of synthetic carbon lines, we have to repeat the same exercise in particular for cooler stars, since Ciialready disappears forTeff>∼30 to 35 kK, in dependence of luminosity class.
From our accumulated experience of analyzing basically all optical carbon lines, we finally pro-vided a list of more than 25 lines of different strength and from different ions, which we consider as reliable carbon abundance indicators.
As a first step toward future work, we studied the direct and indirect (via UV-lines) impact of X-ray emission from wind-embedded shocks onto our sample of optical carbon lines, by simulating an X-ray radiation field that is at the upper limit of realistic values.
In contrast, carbon lines from supergiants with Teff ≈ 35 kK already show changes due to shock radiation, and at 40 to 45 kK the impact of X-rays is strong enough to complicate a reliable abundance measurement. Mostly, the Ciiilines become stronger (corresponding to an increase of up to 0.1 dex in [C/H]; the largest changes occur in the UV-influenced Ciii5696 and Ciii4647-50-51), and the Civ lines become weaker (corresponding to a decrease of 0.3 to 0.6 dex in [C/H]). Consequently, it might become more difficult (or even impossible) to find a value of [C/H] which brings different carbon lines into agreement, when strong X-ray emission would be present. Thus we concluded that in the spectral regime earlier than O7 I, it might be necessary to include the impact of X-rays by default (though a deeper understanding of typical X-ray parameters might be required as well). For supergiants with Teff >
∼ 50 kK, the X-rays lose their impact, since direct ionization due to the stellar radiation field dominates over Auger-ionization and ionization from the EUV-component of the shock radiation (see Carneiro et al. 2016).
This study has been performed as a first step toward a complete CNO analysis. Particularly the investigation of the ratio N/C vs. N/O (as already begun by Martins et al. 2015a,b) will allow us to derive better constraints on the mixing history and chemical enrichment in massive stars than from a pure nitrogen analysis alone. We intend to continue such work, also by including the information provided by UV carbon lines, and by investigating the impact of wind inhomogeneities, which might play an important role also in the UV range, due to porosity effects and because of affecting the ionization balance in the wind.
Table 3.4: Ciilevels: label, electronic configuration, and term designation.
C2 # Configuration Term C2 # Configuration Term 1 1s22s22p 2P 22 1s22s25g 2G 2 1s22s2p2 4P 23 1s22s26s 2S 3 1s22s2p2 2D 24 1s22s2p(3P0)3p 4D 4 1s22s2p2 2S 25 1s22s2p(3P0)3p 2P 5 1s22s2p2 2P 26 1s22s26p 2P0 6 1s22s23s 2S 27 1s22s26d 2D 7 1s22s23p 2P0 28 1s22s26f 2F0 8 1s22p3 4S0 29 1s22s26g 2G 9 1s22s23d 2D 30 1s22s26h 2H0 10 1s22p3 2D0 31 1s22s2p(3P0)3p 4S 11 1s22s24s 2S 32 1s22s27s 2S 12 1s22s24p 2P0 33 1s22s(3P0)3p 4P 13 1s22s2p(3P0)3s 4P0 34 1s22s27p 2P0 14 1s22s24d 2D 35 1s22s27d 2D 15 1s22p3 2P0 36 1s22s27f 2F0 16 1s22s24f 2F0 37 1s22s27g 2G 17 1s22s25s 2S 38 1s22s27h 2H0 18 1s22s25p 2P0 39 1s22s2p(3P0)3p 2D 19 1s22s2p(3P0)3s 2P0 40 1s22s28g 2G 20 1s22s25d 2D 41 1s22s2p(3P0)3d 4F0 21 1s22s25f 2F0
3.A Appendix A: Electronic states of each carbon ion
This section provides a short description of each ion considered in our new carbon model atom, except for Cvwhich is described by the ground level only (see Sect. 3.2.2). All the next three tables have the same format: the first column displays the label of the level, the second column the electronic configuration of that level, and the third column presents the term designation. Table 3.4 refers to Cii, Table 3.5 to Ciii, and Table 3.6 to Civ.
C3 # Configuration Term C3 # Configuration Term
1 1s22s2 1S 36 1s22s5p 3P0
2 1s22s2p 3P0 37 1s22p(2P0)3p 1S
3 1s22s2p 1P0 38 1s22s5d 3D
4 1s22p2 3P 39 1s22s5g 3G
5 1s22p2 1D 40 1s22s5g 1G
6 1s22p2 1S 41 1s22s5d 1D
7 1s22s3s 3S 42 1s22p(2P0)3d 1P0
8 1s22s3s 1S 43 1s22s5f 3F0
9 1s22s3p 1P0 44 1s22s5f 1F0 10 1s22s3p 3P0 45 1s22s6s 3S
11 1s22s3d 3D 46 1s22s6s 1S
12 1s22s3d 1D 47 1s22s6p 3P0
13 1s22p(2P0)3s 3P0 48 1s22s6p 1P0
14 1s22s4s 3S 49 1s22s6d 3D
15 1s22p(2P0)3s 1P0 50 1s22s6g 1G
16 1s22s4s 1S 51 1s22s6g 3G
17 1s22s4p 3P0 52 1s22s6d 1D
18 1s22p(2P0)3p 1P 53 1s22s6h 3H0
19 1s22s4d 3D 54 1s22s6h 1H0
20 1s22s4f 3F0 55 1s22s6f 3F0 21 1s22s4f 1F0 56 1s22s6f 1F0 22 1s22s4p 1P0 57 1s22s7s 3S 23 1s22p(2P0)3p 3D 58 1s22s7p 1P0
24 1s22s4d 1D 59 1s22s7d 3D
25 1s22p(2P0)3p 3S 60 1s22s7g 3G 26 1s22p(2P0)3p 3P 61 1s22s7d 1D 27 1s22p(2P0)3d 1D0 62 1s22s7f 3F0 28 1s22p(2P0)3p 1D 63 1s22s8p 1P0 29 1s22p(2P0)3d 3F0 64 1s22s8d 3D 30 1s22p(2P0)3d 3D0 65 1s22p9d 3D 31 1s22s5s 1S 66 1s22p(2P0)4s 3P0 32 1s22s5s 3S 67 1s22p(2P0)4p 1P 33 1s22p(2P0)3d 3P0 68 1s22p(2P0)4p 3D 34 1s22p(2P0)3d 1F0 69 1s22p(2P0)4p 3P 35 1s22s5p 1P0 70 1s22p(2P0)4p 1D
Table 3.6: As Table 3.4, but for Civ.
C4 # Configuration Term C4 # Configuration Term
1 1s22s 2S 26 1s27i 2I
2 1s22p 2P0 27 1s27h 2H0
3 1s23s 2S 28 1s28s 2S
4 1s23p 2P0 29 1s28p 2P0
5 1s23d 2D 30 1s28d 2D
6 1s24s 2S 31 1s28f 2F0
7 1s24p 2P0 32 1s28g 2G
8 1s24d 2D 33 1s28h 2H0
9 1s24f 2F0 34 1s28i 2I
10 1s25s 2S 35 1s29s 2S
11 1s25p 2P0 36 1s29p 2P0
12 1s25d 2D 37 1s29d 2D
13 1s25f 2F0 38 1s29f 2F0
14 1s25g 2G 39 1s29g 2G
15 1s26s 2S 40 1s29h 2H0
16 1s26p 2P0 41 1s29i 2I
17 1s26d 2D 42 1s210p 2P0
18 1s26f 2F0 43 1s210d 2D
19 1s26g 2G 44 1s211p 2P0
20 1s26h 2H0 45 1s211d 2D
21 1s27s 2S 46 1s212p 2P0
22 1s27p 2P0 47 1s212d 2D
23 1s27d 2D 48 1s213p 2P0
24 1s27f 2F0 49 1s213d 2D
25 1s27g 2G 50 1s214d 2D
This appendix displays the sensitivity of synthetic carbon spectra from dwarf and supergiant models at 30 and 40 kK, with respect to variations inTeff, logg, and ˙M. Figs. 3.16 to 3.19 have the same layout as Fig. 3.6, and refer to Sect. 3.3.3, where also the corresponding discussion has been provided. For completeness, Fig. 3.20 displays the sensitivity with respect to variations in Teff and logg for our complete set of lines, with the same color coding as in Fig. 3.6.
The general behavior is similar to the lines already discussed in this Appendix and Sect. 3.3.3, but not identical, due to different formation regions and processes. (We note that Cii5648 is absent for this model.) The major exception is Cii6461, where the emission displayed by the central model is reduced when any kind of variation inTeff or loggis applied. However, the predicted line is so weak that an actual understanding of the actual origin of this behavior is difficult to obtain.
Figure 3.16: Cii5145, Ciii5696, and Civ5801 line profiles for model D30 (black lines, see Table 3.2) and similar models with relatively small changes in effective temperature (Teff) and gravity (logg). In the upper panels, the red lines correspond to a D30 model withTeff increased by 1000 K, the green lines to a model withTeff decreased by the same value, while the blue lines display the reaction to a decrease of ˙Mby a factor of three. In the lower panels, the red and green lines correspond to a D30 model with loggincreased and decreased by 0.15 dex, respectively
Figure 3.17: As Fig. 3.6, but for model D40, and Ciii4068-70, instead of Cii5145.
Figure 3.18: As Fig. 3.16, but for model S30, and∆logg=0.1.
Figure 3.19: As Fig. 3.17, but for model S40.
Figure 3.20: Carbon line profiles for model D35 (black), and dependence on stellar parameters, for our complete set of lines. The red and green solid line profiles correspond to an increase and decrease by 1500 K inTeff, respectively, whereas the red and green dashed profiles correspond to a variation of +0.2 dex and -0.2 dex in logg.
Surface abundances of CNO in Galactic O-stars:
A pilot study with FASTWIND
This chapter is a copy of Carneiro et al. (2018b), recommended for publication in Astronomy &
Astrophysics by the referee.
Abstract Rotational mixing is known to significantly affect the evolution of massive stars; however, we still lack a consensus regarding the various possible modelling approaches and mixing recipes describing this process. The empirical investigation of surface abundances of carbon, nitrogen, and oxygen (CNO) in large samples of O- and B-type stars will be essential for providing meaningful observational constraints on the different available stellar evolution models.
Setting up and testing adequate tools to perform CNO surface abundance determinations for large samples of O-type stars, by means of the fast performance, NLTE, unified model atmosphere code FASTWIND.
We have developed a set of semi-automatic tools for measuring and analyzing the observed equiv-alent widths of strategic optical C, N, and O lines from different ions. Our analysis strategy is based on a χ2minimization of weighted differences between observed and synthetic equivalent widths, the latter computed from tailored model grids. We pay special attention to the (significant) errors intro-duced by typical uncertainties in stellar parameters. In this pilot study, we describe these tools, and test their performance and reliability using a set of high quality spectra of a sample of 18 presum-ably single Galactic O-type stars with low projected rotational velocities (vsini<∼100 km s−1), and previously determined stellar parameters. In particular, we compare the outcome of our analysis with results from existing studies and theoretical stellar evolution models.
Most of our results for carbon and nitrogen agree, within the errors, with both theoretical expec-tations and literature values. While many cooler dwarfs display C and N abundances close to solar, some of the early- and mid-O dwarfs – and most supergiants – show significant enrichment in N and depletion in C. Our results for oxygen in late-O dwarfs are, however, unexpectedly low, possibly indi-cating deficiencies in the adopted oxygen model atom. For all other objects, no systematic problems
abundances, and we suggest hypotheses regarding their origin.
Our method is (almost) ready to be applied to large samples of late and mid O-type stars – although the oxygen model atom needs to be improved and carefully tested first. For early O-type stars (O4 and hotter), a simultaneous UV analysis seems to be inevitable, due to the scarcity and weakness of optical C and O lines. This will necessarily imply a more complex modeling, additionally accounting for the effects of X-rays from wind-embedded shocks and wind inhomogeneities.