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Chemical compositions of thin-disk, high-metallicity red horizontal-branch field stars

We present a detailed abundance analysis and atmospheric parameters of 76 stars from a survey to identify field Galactic red horizontal-branch (RHB) stars. High-resolution echelle spectra (R\simeq60,000, S/N>=100) were obtained with 2.7 m Smith Telescope at McDonald Observatory. The target stars were selected only by color and parallax information. Overall metallicities and relative abundances of proton-capture elements (C I, N I, O I, Li I), alpha-elements (Ca I and Si I), and neutron-capture elements (Eu II and La II) were determined by either equivalent width or synthetic spectrum analyses. We used CN features at 7995-8040 {\AA} region in order to determine 12^C/13^C ratios of our targets. Investigation of the evolutionary stages, using spectroscopic T_eff and log g values along with derived 12^C/13^C ratios, revealed the presence of 18 probable RHB stars in our sample. We also derived kinematics of the stars with available distance information. Taking into account both the kinematics and probable evolutionary stages, we conclude that our sample contains five thick disk and 13 thin disk RHB stars. Up until now, RHB stars have been considered as members of the thick disk, and were expected to have large space velocities and sub-solar metallicities. However, our sample is dominated by low velocity solar-metallicity RHB stars; their existence cannot be easily explained with standard stellar evolution.


INTRODUCTION
In conjunction with onset of quiescent helium fusion, stars locate on the zero-age horizontal branch (HB) of the HR diagram. Usually HB stars have two energy sources: in addition to the helium burning in their cores, they experience hydrogen fusion in a surrounding shell. The HB domain encompasses a very large effective temperature range, leading to stars being labeled as members of the extreme HB, blue HB, RR Lyrae variables, red HB (RHB), and red clump (RC) stars. The loci of the stars on HB partly depends on many parameters, including stellar mass, metallicity, age, helium abundance, and rotation, and mass. Mass is a key parameter: HB locations depend sensitively on initial stellar masses or the amount of mass loss encountered during the red giant branch (RGB) phase. Theoretical models (e.g. Sweigart & Gross 1978, Piersanti et al. 2004) suggest that very low mass stars (around 0.5 M ⊙ ) appear on the blue part of the HB, higher-mass ones (around 0.8 M ⊙ ) sit on the RHB, and stars of solar mass and greater are located in the small RC area at the extreme red edge of the HB.
RHBs are located between the instability strip and the RC. Stars of the RC concentrate at M V ≈ 1 and B −V ≈ 0.9-1 (Cannon 1970, Wilson 1976. These colors correspond to effective temperatures T eff ≈ 5000-4800 K (e.g, Ramírez & Meléndez 2005), in contrast to the color/T eff range of RHB stars (B −V ) 0 ≈ 0.5-0.8, T eff ≈ 5000-6200 K (Straizys et al. 1981, Gray & Corbally 2009. It is easy to identify RHB stars in globular cluster (GC) color-magnitude diagrams (e.g, see the survey of Rosenberg et al. 2000a, Rosenberg et al. 2000b) but they are not easily distinguished from RC, SGB (subgiant branch), and even MS (main sequence) stars in the Galactic field. There have been several attempts to isolate field RHB stars. An early systematic survey was that of Straizys et al. (1981), who observed eight RHB candidates with the Vilnius photometric system. They derived the metallicities and determined the luminosity classes using color-color diagrams, but could not prove their RHB membership because of lack of surface gravity information.
Field RHB stars were first investigated as a group by Rose (1985). He studied G5-G7 stars selected from the Upgren (1962Upgren ( , 1963 North Galactic Pole (NGP) survey. By analyzing low resolution spectroscopic data, Rose showed that one can distinguish between evolved and "post-main sequence stars" (i.e, subgiants) using the strengths of Sr II λ4077 A and CN λ3883 and 4216Å bands. He also estimated the scale height of RHBs to be > 500 pc and concluded that they are moderately metal-poor members of the "thick disk" (Gilmore & Reid 1983) of the Galaxy.
The evolutionary states of alleged RHB field stars were revisited by Tautvaišienė (1996), who carried out Vilnius system photometric observations of 13 field RHB candidates from Rose (1985). From these data she deduced spectral types, atmospheric parameters and absolute magnitudes. Her results yielded a complex message: low mean metallicity (<[Fe/H]> ≈ −0.6) and large mean age (t ≈ 10-12 Gyr) consistent with possible RHB status, but low mean temperature (<T eff > ≈ 4940 K) consistent with the RC.
In another study, Tautvaišienė (1997) gathered high-dispersion spectra of 10 field RHB stars that had been identified in a number of previous papers. The metallicity range of her sample was very large, −0.2 ≥ [Fe/H] ≥ −1.9, which resulted in mixed Galactic population membership. She concluded that the chemical compositions of these RHB stars were in accord with those of dwarf and red giant stars of similar metallicities, and thus RHB stars can be used as tracers of the chemical evolution of the Galaxy. Tautvaišienė et al. (2001) followed this paper with another chemical composition study of Rose (1985) stars, finding in particular: (a) overabundances of α-elements (O, Mg, Si, Ca, Ti); (b) solar abundances of some neutron-capture elements (Y, Ba, La) and overabundances of others (especially Eu, possibly Zr and Sm): (c) depleted C and enriched N abundances, indicative of CN-cycle H-fusion in these stars; and (d) perhaps most importantly, very low carbon isotope ratios, 12 C/ 13 C ≈ 3-6.
Previous studies have suggested that true field RHB stars are relatively rare, but the extant sample sizes are small. The ones that exist should be relatively low mass members of the thick disk because stellar evolution models (e.g. Lee et al. 1994) predict that higher-mass, solar-metallicity core He-burning stars should reside exclusively in the RC. Following the Norris (1987) suggestion that, "the only way to resolve this problem in an unambiguous way will be to analyze a complete sample of stars having colors in the range 0.7 ≤ B −V ≤ 0.95", Kaempf et al. (2005) identified a large sample of RHB candidates from the Hipparcos catalog (Perryman et al. 1997). Their kinematic analysis found both thick disk and halo RHBs, and they estimated a Galactic scale height of 0.6 kpc for the thick disk RHBs.
In this paper we report on a new large-sample chemical composition study of candidate field RHB stars. The combined kinematics, metallicities, and chemical abundance ratios will be used to confirm the existence of a substantial number of high-metallicity thin-disk RHB stars. In §2 we present the target star selection and the spectroscopic data that were gathered. Model atmosphere parameter derivation is described in §3. We discuss the chemical composition analysis in §4, along with comparison to other large-sample literature studies. The kinematics of our sample are given in §5, and estimates of their evolutionary states in §6. We discuss the implications of these results in §7.

OBSERVATIONS
We obtained high resolution, high signal-to-noise spectra of stars selected mainly by their colors (as suggested above) and some indications that they might be giants (spectral luminosity class III and/or low M V from Hipparcos parallax).

Selection of the Program Stars
We assembled our program star list from several sources. Most of our RHB candidates were selected from among the Upgren (1962) NGP survey and the Harlan (1969Harlan ( , 1974Harlan ( , 1981 stars used by Rose (1985). We also made use of SIMBAD 1 and Vizier 2 .
We restricted the prospective sample in color, spectral type, and apparent magnitude (V < 11). For spectral type, we limited our search to between G0III and G8III. The reason for this restriction is simple: the loci of the RHB stars in the HR diagram fall in between the instability strip on the blue side and RC region on the red side. RHB stars have very similar luminosities to RC stars, although the effective temperatures of RCs are lower. For RC stars Kaempf et al. (2005) adopted a luminosity-dependent color range of 0.85 B −V 1.2 (spectral classes approximately G6-K3), and for RHB stars 0.5 B −V 0.8 (approximately G0-G5); see their Figure 1. In our color selection we used this RHB B −V range, but when possible we preferred to use an equivalent V − K range of ≈1.5-2.2, because V − K colors are almost independent of metallicity and gravity but very sensitive to effective temperature. In order to calculate photometric temperatures we have used the metallicity-dependent T effcolor formula for giants given by Ramírez & Meléndez (2005). The apparent magnitude limit was defined by our telescope/instrument combination; see §2.2. The radial velocity and proper motion properties of the candidates were not included in the selection criteria, which allowed us to minimize potential kinematic biases in our sample. There is an inevitable -5bias due to the cutoff in apparent magnitude limit of the instrumental set-up.

Observations and data reduction
We obtained high-resolution spectra of RHB candidates with the Robert G. Tull Cross-Dispersed Echelle spectrograph (Tull et al. 1995) of the 2.7m Harlan J. Smith Telescope at McDonald Observatory. The instrumental setup, including a 1. ′′ 2 width entrance slit, yielded a spectral resolving power of R ≡ λ/∆λ ≈ 60000. Within the wavelength coverage of λλ 3400-10900, we observed 63 spectral orders. The gaps between the orders reduce towards blue orders, enabling continuous spectral coverage for λ 5900Å. The data were obtained during five observing runs corresponding to 22 nights in total.
Based on the selection criteria discussed in §2.1, we were able to observe 129 candidate RHB stars with adequate signal-to-noise. Not all of the stars survived scrutiny after data reduction. Some of these stars were of interest because they were being observed with high spectral resolution for the first time. Inspection of the data revealed that 20 candidates either were rapidly rotating (V sin i ≥ 30) or were double-lined spectroscopic binaries. These were discarded from the final sample because in both cases the spectral features become too blended and do not allow us to extract meaningful abundances for this project. In the end we also had to eliminate 33 more candidates during the model atmosphere analyses because we were not be able to obtain reasonable atmospheric parameters from spectroscopic criteria alone (see §3 for further discussion). The remaining sample of 76 candidates were studied in detail. In Table 1 we present basic data for these program stars.
The data reductions were performed with the IRAF 3 software package in the classical manner: bias subtraction, flat-field division, and scattered light removal followed by extraction of the spectral orders. We took at least two exposures per star and combined them in order to filter out cosmic rays events. Fast-rotating hot stars were observed each night at appropriate air-masses, and used to remove telluric features from the spectra of the program stars. Removal of the telluric lines was done with the IRAF routine telluric, following continuum normalization. ThAr lamp exposures taken at the beginning and end of each night were used for the wavelength calibration, again done with standard IRAF tasks.
Since some of the RHB candidates have reported parallaxes but no radial velocities (RVs), we also observed a few IAU RV standards. We used these standards to measure the RVs of our stars via the cross-correlating technique provided in fxcor (Fitzpatrick 1993) task in IRAF. The mean RVs of the relevant stars are listed in Table 1. Typical errors for the RV measurements were around ±0.3-0.6 km s −1 . We also checked our measurements by re-measuring the RVs of a few of our program stars that have previously published values. Our RVs agree with the literature values to within the joint error estimates.
We estimated signal-to-noise (S/N) ratios of the spectra at three different "continuum" wavelength regions near 4500, 5500 and 6500Å in which we detect no obvious absorption features. We give examples of the derived S/N for three program stars with different visual magnitudes in Table 2.

MODEL ATMOSPHERE PARAMETERS
We have used both equivalent widths (EW s) and spectrum synthesis for the abundance analysis. The line list for model atmospheric parameter determination was generated by choosing clean lines with laboratory gf -values between λ5500-6500Å. This spectral region contains many relatively unblended Fe-peak species transitions and only very weak CN molecular contamination in our program stars. In selecting the transitions we consulted the solar line compendium of Moore et al. (1966) and the Griffin (1968) Arcturus spectral atlas. We measured EW s using the SPECTRE 4 code (Fitzpatrick & Sneden 1987). SPECTRE employs a semi-automated routine that fits Voigt and Gaussian line profiles to the observed spectral lines.
In Figure 1 we compare our EW measurements with the values given by Takeda et al. (2005) for HIP 13339 and HIP 71837. The EW scales are in reasonable agreement: defining ∆EW ≡ EW literature − EW thisstudy , we found <∆EW> = +1.6 mÅ, with σ ≃ 5 mÅ. Stellar atmosphere models required by the analysis were from the Castelli et al. (1997), Castelli & Kurucz (2003) grid 5 of model atmospheres computed with opacity distribution functions and without convective overshooting. We interpolated in this grid with software developed by Andy McWilliam and Inese Ivans. Then we used an automated version of the spectral line analysis and synthetic spectrum code MOOG (Sneden 1973) 6 to determine the abundances of our program stars. MOOG performs analyses using one-dimensional 4 An interactive spectrum measurement package, available at http://www.as.utexas.edu/∼chris/spectre.html 5 Available at: http://kurucz.harvard.edu/grids.html 6 available athttp://www.as.utexas.edu/∼chris/moog.html local thermodynamic equilibrium (LTE) equations for plane parallel atmospheres. In the automated version (described in more detail by Hollek et al. 2011 and Roederer et al., in preparation) all of the steps in model atmosphere derivation and many parts in the relative abundance analysis are done iteratively by the code without human intervention.
We used Fe I and Fe II abundances derived with trial model stellar atmospheres to determine the fundamental stellar parameters of effective temperature T eff , surface gravity, log g, microturbulent velocity ξ t , and metallicity [Fe/H] 7 . Final values of T eff were estimated by requiring that Fe I abundances show no trend with excitation potential χ beyond internal line-to-line scatter uncertainties. A similar approach yielded estimates of ξ t , which is related to small-scale turbulent motion: this parameter was varied until we obtained no apparent Fe I abundance trend with reduced width RW (≡ EW/λ).
As reported in previous studies, non-local thermodynamical equilibrium (NLTE) mechanism almost has no effect on Fe II lines and has a small effect on Fe I lines for solar-type metallicities and mildly metal-poor stars. As given by Mashonkina et al. (2011) for a temperature range of 4600−6500 K, the departure from LTE is not more than 0.1 dex in Fe I abundances. This value is well within the uncertainty limits of our Fe I abundances (see §3.1). Since we have mostly solar-metallicity and mildly metal poor stars in our program, the corrections for NLTE effects were not applied.
Surface gravities were calculated by demanding that Fe I and Fe II lines yield the same mean abundances to within the 1σ internal scatter uncertainties of both species. Finally, we iterated on the [Fe/H] metallicities until the values assumed in model creation were consistent with those implied by the line abundance averages. Derived model parameters of our stars are given in Table 3.
We also derived solar abundances (Table 4) applying the same procedure to all the species investigated here. The high-resolution solar data were obtained from the electronic version of the solar center-of-disk spectrum of Delbouille et al. (1973) 8 . These abundances were used for differential determination of the stellar abundances. These differential abundances should give more accurate "internal" results than absolute abundances since many of the systematic errors can be nearly compensated in comparing our stars with the Sun. Our internal solar abundances generally agree with those recommended by Asplund et al. (2009), to within the mutual uncertainties. Nitrogen is the sole exception with a substantial difference, ∼+0.15 dex, which potentially could lead to a small offset in our [N/Fe] results compared to the previous studies that we will consider in §4.3. Our adopted CN oscillator strengths from the Kurucz database may be the major factor contributing to this offset. Exploration of this issue in detail is beyond the scope of our paper.

Uncertainties in Parameters
The internal uncertainties in atmospheric parameters T eff , log g, and ξ t were estimated by running a series of trial analyses on the spectral data of BD+27 2057, HD 84686 and HIP 98587. For temperature uncertainties, we varied the assumed T eff in steps of 50 K and kept the other parameters fixed during the analysis. The temperature was changed until the mean abundance difference between low and high excitation Fe I lines exceeded the ±1σ scatter of individual line abundances derived with the optimal T eff , in other words, until Fe I abundances showed an obviously unacceptable trend with excitation potential. This method yielded an average uncertainty of ∼150 K for the T eff .
To estimate the external uncertainty in T eff we have searched the literature for previous high-resolution studies of our program stars. We found no large samples in common with our stars, but from the T eff values reported in various publications we have generated Figure 2. The heterogeneity of the literature data does not justify a detailed statistical treatment, but it is clear that our T eff values track those in previous publications, with a scatter of ∼120 K. We adopt an overall T eff uncertainty of ±150 K.
Our T eff values for a few stars, however, deviate significantly from previous estimates. Since temperature estimates are critical to assessment of RHB or RC status for our stars, we turned to the spectroscopic "Line Depth Ratio" (LDR) method of T eff determinations. Gray & Brown (2001) showed that simple LDRs of several line pairs in the 6200Å spectral region (most often a ratio formed by comparing the central depths of a V I and an Fe I line) Their paper maps the LDRs to stellar B − V values and those colors in turn to T eff estimates. LDRs are excellent temperature indicators especially for giants with spectral types between G3 and K3 (see also Figure 7 of Gray & Brown 2001). This method produces effective temperatures well within the uncertainty levels of 150 K and gives us an independent spectroscopic check on our spectroscopic T eff estimates from Fe I lines.
We measured LDRs for our whole sample. We mainly made use of the line ratio pairs recommended by Gray & Brown (2001) 9 and derived an average T eff (LDR) for the stars. These line pairs compare the central depths of a very T eff -sensitive transition (almost always a low-excitationV I line) with usually a high-excitation transition of another Fe-group species. We especially made use of the 6224.5Å/6223.9Å pair that compares a V I line with χ = 0.29 eV to a Ni I line with χ =4.10 eV. Their LDR was the most useful because it could be reliably measured throughout most of the T eff domain of our sample. However, the depths of V I lines weaken rapidly with increasing temperature (≥ 5500 K). For example, for the five V I lines recommended by Gray & Brown (2001) for LDR studies, < EW > ≃ 6 mÅ in the solar spectrum (Moore et al. 1966). Lines with such small EW s are near the reliable detection/measurement limit of our spectra. Therefore for warmer stars of our sample we also measured some LDRs proposed by Strassmeier & Schordan (2000).
A comparison of our spectroscopic and LDR temperatures are given in Figure 3. The error bar depicted in this figure corresponds to the standard deviation of the mean difference between these two temperatures, σ ≃ ±150 K. A change in the correlation is apparent at T eff < 5500 K, which we have highlighted by using different colors for the symbols of stars warmer and cooler than this T eff . Although we see a small systematic bias towards higher temperatures for T eff < 5500 K (Figure 3, purple points), the standard deviation of these points is around 90 K and well within the uncertainty limits. We suggest that this systematic deviation may both arise from the method we adopt from Gray & Brown (2001), which first involves a calibration of LDR against B−V color indices, then converts these into temperatures, and LDRs that we obtained with different instrumentation than those of Gray & Brown (2001). As mentioned above, this temperature is near the reliability limit for the LDR method using these transitions. On the other hand, for T eff < 5500 K, which proves to be the more important domain for identifying true RHB targets, these two T eff estimators are considerably in good agreement. This result suggests that for stars when clashes occur between photometric and spectroscopic temperatures, the photometric values are probably not reliable for T eff < 5500 K. In general, LDR method supports our claimed temperatures for the stars, including the ones which our spectroscopic values clash with literature estimates.
We determined internal uncertainties for log g and ξ t through repeated trials with variations in these quantities. The typical average uncertainties were estimated as 0.16 dex and 0.2 km s −1 , respectively. Comparison of our log g results with previously reported ones resulted in an external uncertainty level of ≈0.25 dex. By taking into account both internal and external uncertainty levels, we adopt an average uncertainty for log g of ≈±0.3 dex.
We also calculated "physical" gravities and compared them with spectroscopic gravities in order to possibly gain insight on the masses of our stars. We use the following standard V I(6256.8Å)/Fe I(6255.9Å) equation for physical gravities For the Sun, we adopted T eff = 5780 K, log g= 4.44 dex, and M Bol⊙ = 4.75 mag. Bolometric corrections were calculated using the relation given by Alonso et al. (1999). A comparison of physical gravities, log g(theo.), with spectroscopic gravities, log g(spec.), is given in Figure 4. A significant uncertainty in a physical gravity calculation is the assumed mass. Therefore in the figure we give three log g(theo.) values calculated for the masses of 1, 2 and 4 M ⊙ . They are shown with dashed and dotted lines. The best average agreement between these two log g scales for significantly evolved stars (log g 3.0) is for M ≃ 2 M ⊙ and for MS/SG stars it is M ≃ 2.5 M ⊙ . The joint uncertainties are too large to draw conclusions about individual stars, but does suggest that our sample is not dominated by high mass stars.
In order to derive the effect of these uncertainties on abundance determinations, we did multiple analyses by changing T eff , log g, [Fe/H], and ξ t within their uncertainty limits. The effective temperature uncertainties create an uncertainty around 0.15-0.16 dex in Fe I and the uncertainties in log g make the Fe II abundances change in less than 0.1 dex. The other elements' abundances vary in between 0.1-0.16 dex due to T eff uncertainties. In general, the uncertainties in surface gravity and microturbulent velocity result in abundance uncertainties less than ≃0.1 dex.
Finally we comment on the uncertainties in our derived 12 C/ 13 C ratios. These are nearly insensitive to model atmosphere uncertainties, since 12 CN and 13 CN are nearly identical molecules. We estimated the 12 C/ 13 C uncertainties by fitting synthetic spectrum to the observed 12 CN and 13 CN features with various isotopic ratios to estimate their maximum and minimum probable values. The error estimates for the 12 C/ 13 C ratios are given in Table 5.

Photometric Temperatures and Reddening
Since we observed almost half of our program stars at high spectral resolution for the first time, our sample lacked a consistent set of reported atmospheric parameters. This made it difficult to have meaningful temperature estimates to start our analyses, so we relied exclusively on our Fe I line analyses to derive T eff values, augmenting these with LDR estimates. Armed with these results we now compare spectroscopic and photometric temperatures for the stars in common. We concentrate on the V − K color because it is one of the best photometric temperature indicators, having almost no dependence on metallicity and gravity (e.g. Alonso et al. 1999, Ramírez & Meléndez 2005. The V magnitudes that we employ are on the standard Johnson photometric system; the K s magnitudes are from the Two Micron All Sky Survey. We adopt the (V − K s )-T eff calibration equation from Ramírez & Meléndez (2005).
Before calculating the photometric T eff values, we first applied interstellar reddening corrections to the V − K s colors. A recent survey of the regions within 300 pc of the Solar System showed that the "Local Bubble" (Lallement et al. 2003) has a shape of an irregular lacuna, which extends ∼60 pc towards the Galactic center, 80−150 pc towards the outermost edge of the Galaxy and ∼200 pc above and below the Galactic plane. But ISM dust extinction is patchy and varies considerably toward different sightlines, and many of our stars have poorly-constrained distances. Therefore we estimated their reddening in three different ways. In all cases, we adopted a typical Local Bubble radius of 75 pc and assumed no reddening for stars estimated to be within that distance (e.g. Henry et al. 2000).
For our first reddening estimate, we have applied an isotropic reddening correction for the stars have distances 75 < d ≤ 300 pc, adopting A v = 0.8 mag kpc −1 (Henry et al.).
The extinction values of eight stars with distances d > 300 pc were obtained from the NASA/IPAC Extragalactic Database (NED) 10 extinction calculator. Zero reddening was assumed for the stars with high Galactic latitudes, |b| ≥ 50 o (e.g, Sandage 1972). For the second and third reddening estimates, we employed the methods described by Chen et al. (1998) and Hakkila et al. (1997). Our computations with the latter method used Hakkila's EXTINCT code 11 which takes into account various interstellar extinction correction methods from several studies (see Hakkila et al. for more details).
We applied each of these reddening estimates to the observed V − K colors, computed photometric temperatures from the reddening-corrected V −K colors using the Ramírez & Meléndez (2005) formula, and correlated them with our spectroscopic temperatures as shown in Figure 5. Inspection of this figure clearly suggests that the two T eff scales are well correlated, and that various reddening assumptions produce little substantive variations in the temperature correlations. Therefore, we have adopted our E(B −V ) values for reddening for the entire sample.

ABUNDANCE DETERMINATIONS
Using the model atmospheres given in Table 3, we determined abundances of several elements in our program stars. We report these as relative abundance ratios [X/H] in Table 7. Whenever possible, we computed the abundances from line EWs. For complex transitions, those that have significant hyperfine and isotopic substructure, and those that have significant line blending issues, we resorted to synthetic/observed spectrum matches to determine the abundances. We were especially interested in three element groups that could help constrain the Galactic population memberships of our stars and verify their evolutionary states. In the following subsections we consider in turn the α, the neutron-capture, and the proton-capture element groups.

α-Elements: Silicon and Calcium
The abundances of α-elements exhibit the same behavior in giant and dwarf stars of solar metallicity ([α/Fe] ∼ 0; Soubiran & Girard 2005, Mishenina et al. 2006. In thick disk main-sequence and subgiant stars, the α elements typically become overabundant as metallicity decreases, reaching [α/Fe] ∼ +0.3 at [Fe/H] ∼ −1 (e.g. Tautvaišienė et al. 2001, Tautvaišienė et al. 2001, Reddy et al. 2006. The temperatures and gravities of these prior large samples are 5000 K T eff 6500 K and 3.4 log g 4.7. Less systematic abundance trend information is available for thin/thick disk giant stars, which in our case cover the approximate parameter range 4800 K T eff 5600 K and 2.2 log g 4.0. The easily-observable α elements are Mg, Si, and Ca. Often Ti is grouped with the other α's due to its similar abundance behavior with metallicity. However, Ti is not a pure α element because its dominant isotope is 48 Ti 22 , which is not an even multiple of α particles. Here we concentrated on Si and Ca abundances, since they have large numbers of transitions with a range of line strengths in the yellow-red spectral region (most Mg I lines are very strong in our G-K giant stars). We generally used about 9 Si I and 15 Ca I lines, and derived abundances from their EWs.
We compare our Si and Ca abundances with published values in Figure 6. The literature data are taken from several studies that are given in the figure caption: Tautvaišienė et al. (2001), Reddy et al. (2003), Bensby et al. (2003), Reddy et al. (2006) and Mishenina et al. (2006). Our sample has a wide metallicity range −1.0 ≤ [Fe/H] < +0.5 and it is clear that the α elements behave similarly in our cooler giants as they do in warmer main sequence stars and subgiants. Even though slight scaling differences (< 0.1 dex) naturally result from different solar abundances, different line choices, and different oscillator strengths adopted

Neutron-Capture Elements
Among thin and thick disk stars, those elements whose solar-system origin is due chiefly to "slow" neutron bombardment reactions (the s-process; e.g. Sr, Y, Zr, Ba, La, and Ce) generally exhibit their solar abundance ratios ([X/Fe] ∼ 0) throughout metallicity regime −1.0 [Fe/H] +0.2. But those elements that are products of "rapid"-blast neutron capture events (the r-process; Eu, Gd, Dy) tend to increase in relative abundance with decreasing metallicity. The most easily-observed r-process element is Eu, and its mean abundances reaches [Eu/Fe] ∼ +0.4 at [Fe/H] ∼ −1. A good summary of the disk s-and r-process abundance trends can be seen in Figure 17 of Reddy et al. (2006). Additionally, Simmerer et al. (2004) has shown that a distinct kinematic signature in disk neutron-capture elements: stars with larger space motions have often much lower s-process abundances than r-process ones (see their Figure 12).
There are not a lot of useful neutron-capture element (Z > 30) transitions in the yellowred spectral region. We considered just La (75% s-process origin in solar-system material, e.g., Sneden et al. 2008 and references therein), and Eu (97% r-process).
La II: transitions at 6262.2 and 6390.5Å can be detected in nearly all of our stars. These lines have well-determined transition probabilities and hyperfine structure parameters (Lawler et al. 2001a; see also Ivans et al. 2006 for complete substructure line lists for these transitions). The sole naturally-occurring isotope of this element is 139 La. Because of the complexity of these La II transitions we used synthetic spectrum analyses to determine the abundances. Notable contaminants to the La features are CN red-system lines, but in most cases the CN strengths were small and did not materially affect the derived abundances.
Eu II: transitions at 6645.1 and 7217.5Å were used in the Eu abundance analyses. The transitions are very complex, because Eu has both hyperfine and isotopic substructure. There are two Eu stable isotopes, 151 Eu (48% of the total solar-system Eu abundance) and 153 Eu (52%). There is little difference in Eu isotopic ratios generated in r-process and sprocess environments. Spectroscopic studies of Eu II in extremely metal-poor, Eu-enhanced stars suggest that the 151 Eu fraction is ≈50% for r-rich stars (Sneden et al. 2002, Aoki et al. 2003a, Roederer et al. 2008) and ≈57% for s-rich stars (Aoki et al. 2003b). Therefore in our Eu syntheses we adopted the solar-system isotopic ratios. The basic laboratory analyses of these lines were published by Lawler et al. (2001b); full hyperfine/isotopic substructure lists are in Ivans et al. (2006).
The abundance variations for La II and Eu II as a function of metallicity are given in Figure 7. The observed trends (no apparent change in [La/Fe], increasing [Eu/Fe] at lower metallicities) are in complete agreement with previous results cited above. The neutroncapture element abundances, combined with those of the α elements clearly provide Galactic population indicators for our stars.

The Proton-Capture Abundances
Evolved Population I solar-metallicity stars show evidence of convective envelope mixing; their observed light elements Li, C, N, and O clearly have been altered via interior synthesis in p − p and CNO cycles. The original C abundance drops by a factor of about two, the N abundance rises by comparable amounts, and the carbon isotopic ratio drops to values usually between 15−30 (e.g. Lambert & Ries 1981). Mildly metal-poor, high-velocity stars (thick disk, labeled old disk in early papers) of about a solar mass have lower 12 C/ 13 C ratios but less evidence for depleted C and enhanced N (e.g., Cottrell & Sneden 1986). Low mass giant stars of the Galactic halo often display more dramatic 12 C→ 13 C and C→N conversions (e.g., Gratton et al. 2000). Thus the CNO abundances of evolved stars can yield information on their internal evolutions and population memberships.
Since stars easily destroy Li in relatively low-temperature proton-capture reactions, Li abundances or upper limit estimates in metal-rich stars provide additional information. The presence of Li in the spectrum of an evolving star usually suggests that envelope convection has not yet developed to a point where dredge-down of original surface Li has effectively cleaned it from the star's envelope. The absence of Li can indicate (but not always) that other proton-capture products may have been dredged up to the surface. This straightforward interpretation is complicated by the relatively rare phenomenon of Li-rich giants (e.g., Charbonnel &Balachandran 2000, Kumar et al. 2011 andreferences therein). This increases the importance of searching for this element in the spectrum of any evolved star.
Unfortunately, many atomic and molecular transitions that are CNOLi abundance indicators have significant detection and/or analytical issues. Here we describe the transitions that we used, and our atomic/molecular parameter choices.
C I: High excitation (χ 7.7 eV) lines of this species are strong in the warmest candidate stars but weaken with decreasing T eff . We adopted the transition probabilities for these lines recommended in the NIST Atomic Spectra Database (Ralchenko et al. 2011) 12 . However, the C I lines decrease in strength so quickly with decreasing T eff , often only a couple of them could be used for C abundances in our stars.
CH G-Band: We also employed the CH X 2 Π−A 2 ∆ "G-band" to derive C abundances via spectrum syntheses. Significant G-band absorption occurs in the 4200−4400Å spectral range, and we estimated C abundances from three regions (4300−4308Å, 4308−4315Å, and 4322−4327Å) and averaged the results. The 4308−4315Å region proved to be the most reliable C abundance indicator, due to greater atomic-line contamination in the other two spectral regions. The synthetic spectrum line list was formed with 12 CH and 13 CH lines from Plez (private communication; see, e.g., Hill et al. 2002, Plez & Cohen 2005 and atomic lines from the Kurucz (2011) compendium. The CH lines are plentiful and relatively strong throughout the atmospheric parameter domain of our program stars. As an example of the CH observed/synthetic spectrum matches, we show in Figure 8 the 4308−4315Å spectral region in a typical program star. Generally more reliable C abundances were obtained from the CH band syntheses than those from the C I EW measurements, especially among the cooler (T eff < 5500 K) stars.

CN:
We used synthetic spectrum calculations of 12 CN and 13 CN A 2 Π − X 2 Σ red system lines in the 7995−8040Å region to determine N abundances and 12 C/ 13 C ratios. The synthesis line lists were taken from Kurucz (2011). A triplet of 13 CN lines near 8004.7Å was the primary 12 C/ 13 C indicator, but several other, usually weaker features (e.g., at 8007.9, 8010.4 and 8011.2Å) were used for confirmation. In Figure 9, we give an example of a synthetic spectrum fit to the 12 CN and 13 CN features in the 8002.5−8011.5Å region. Derived 12 C/ 13 C ratios for our program stars are given in Table 5. C-N-O abundances are bound to each other through molecular equilibrium, which becomes more important towards lower temperatures. The accuracy of N abundances especially depend on the accuracy of C abundances through CN formation. Note that our line list produced a systematic offset of ∼0.15 in the solar N abundance (Table 4) compared to Asplund et al. (2009). This would suggest that perhaps we have an offset of about +0.15 dex in our N abundances.
[O I]: There are two ground-state forbidden transitions upon which are based most of the O abundances in red giant stars. We analyzed only the 6300.3Å; its companion at 6363.8Å is much weaker and suffers from large amounts of CN contamination. The 6300Å line has a very accurate transition probability (Allende Prieto et al. 2001 Johansson et al. (2003).
O I: In the spectra of our warmer stars the forbidden line is too weak to yield reliable abundances, so we also used the often-analyzed very high-excitation 7770Å triplet. For these lines we adopted gf -values from the NIST database (Ralchenko et al. 2011). Unfortunately, it is well known that the O I triplet lines are subject to NLTE effects, in the sense that their LTE-based abundances are always too large ( 0.2 dex, Gratton et al. 1999, Bensby et al. 2004) compared to abundances from the forbidden lines. In order to correct for the NLTE effects, we followed a similar approach to the NLTE-correction method of Bensby et al. (2004). That is, we applied a "robust regression" analysis and derived the following equation  Figure 10 In Figure 11, we plot our CNO abundances along with previously reported values, using the same comparison samples. While the derived mean O and N abundances of our stars are are in good agreement with literature mean values at similar metallicities, our C abundances are somewhat lower those from previous studies. But the comparison samples in this case are not the best, because the literature data are dominated by main sequence and subgiant stars, while our program stars are clearly much more evolved. Therefore in Figure 12 we repeat the exercise of the previous figure, but with two differences. First, we have eliminated the stars of our sample that are probably not chemically mixed (those with T 5400 K and/or log g 3.5). Second, we added literature CNO abundance results only from studies that concentrated on red giant stars. Our C abundances are more in accord with those in other samples of evolved giants.
Thick disk giants ([Fe/H]<-0.25) in general show higher C and O abundances than their thin disk counterparts. The O enrichment is indicative of larger Type II SNe contributions to thick disk than thin disk stars, combined with no depletion in ON-cycle H-fusion (interior temperatures are too low) in the low-mass thick-disk giants that we observe. The higher C in thick disk stars suggests that the lower-metallicity, relatively low-mass thick-disk giants have relatively shallow convective envelopes, and dredge-up has failed to reach the interior fusion zones where the CN-cycle has run to completion. This echos the discussion in Cottrell & Sneden 1986, who concluded that in their old disk giants, "only the coolest (outer) portions of the CNO-processed hydrogen burning shell, where 12 C(p, γ) 13 C has taken place, have been convectively mixed into outer layers". Theoretical support for this notion was given in Sneden et al. (1986), whose Table 6 gives D. A. VandenBerg's predicted C/N ratios from standard evolutionary model stars that have experienced first dredge-up. It is clear that as as mass and metallicity decrease, model stars exhibit decreasing surface C and N abundance changes.
Li I: Unless the Li abundance is extraordinarily large, only the 6707.8Å resonance line of the neutral species can be detected in cool stars. This ground-state doublet transition has a well-determined gf value; see the summary of experimental and theoretical studies in Smith et al. (1998). The transition is complex, with both isotopic and hyperfine substructure. We adopted the parameters given in the line compendium of Kurucz (2011) 13 .
We were able to measure Li I abundances for some of our targets (see Table 6). Since Li line strength is a severe function of T eff , for stars with undetectable Li here are approximate abundance upper limits as guidelines for interpretation: for stars with log g ≥ 3.5 and T ≥ 5500, logǫ(Li) < 1, and for those with log g < 3.0 and T ≤ 5100, logǫ(Li) < +0.5 (with the limit of course decreasing towards lower temperatures). The Li abundances decrease as T eff and log g decrease. This is as expected in normal stellar evolution: during the path from main sequence to red giant phases, the deepening convective envelopes bring Li down from the surface layers where it can be destroyed easily at relatively low fusion-zone temperatures. Program stars with higher surface Li abundances have undetectable 13 C, consistent with their chemically un-evolved status. Six stars of our sample have low 12 C/ 13 C values, but detectable Li I lines; these all appear to be plausible RHB and RC candidates (see Table 7). For these stars, <log ǫ(Li)> ∼ 0.8. Standard evolution models (e.g., those without extra mixing mechanisms seeming needed for Population II giants) predict surface Li depletion by factors up to about 30 (Iben 1967a). If a star begins main-sequence life with the present interstellar medium Li abundance (log ǫ(Li) ∼ +3.0, e.g. Grevesse & Sauval 1998) then through ordinary stellar evolution it should exhibit log ǫ(Li) +1.5. None of our evolved program stars have Li abundances near to this value, indicating either that their natal Li contents were not as high as the present ISM value, or that extra mixing during their lives destroyed Li beyond the standard predictions.

KINEMATICS
Among the overall sample of 76 stars, we have distance information for 58 of them to facilitate a kinematical investigation. In order to calculate stellar space velocities, we have collected parallaxes and proper motions, as well as their errors, from van Leeuwen (2007), who recently published new reduction results of HIPPARCOS data. We either adopted the RVs already available in the literature or we measured new values as described in §2.2. We used the matrix equations given by Johnson & Soderblom (1987) and calculated the Galactic velocity components U LSR (positive toward the Galactic center, V LSR (positive in the direction of the Galactic rotation), W LSR (positive toward the NGP), and their uncertainties (σ U LSR , σ V LSR , σ W LSR ) with respect to the Local Standard of Rest (LSR). Correction for the solar motion was made by adopting the values (U, V, W ) ⊙ = (+10.00, +5.25, +7.17) from Dehnen & Binney (1998). The results are given in Table 8.
In Figure 13 we show these kinematics in a Toomre diagram. This kind of plot gives the kinematical distribution of the stars in terms of their combined vertical and radial kinetic energies as a function of their rotational energies. The solid curves of Figure 13 are lines of constant total space velocity, V tot = (U 2 LSR + V 2 LSR + W 2 LSR ) 1/2 . The dashed line represents the V tot = 70 km s −1 , which approximately separates the thin and thick disk populations of the Galaxy (e.g. Bensby & Feltzing 2009, Nissen 2004. Stars with total space velocities of 70 < V tot < 180 km s −1 are considered to be probable thick disk, those in the range 50 < V tot < 70 km s −1 are called transition (thin/thick) objects, and those with V tot < 50 km s −1 are probable thin disk members of the Galaxy. Figure 13 reveals that our sample is dominated by thin disk stars. Out of 58 stars with kinematic information, 39 (the squares) probably reside in the thin disk, 10 probably are in the thick disk, and 9 are kinematic transition objects whose thin/thick status is ambiguous. In Figure 13 we also use filled and open version of the symbols to indicate the detection status of 13 CN features in individual stars: those with 12 C/ 13 C < 30, and those with no detection of 13 CN (thus 12 C/ 13 C > 30), respectively. We will consider the carbon isotopic ratio distribution further in §7.

Inspection of
Looking at the kinematics of our stars in more detail, in Figure 14 we show the correlations between total space velocity V tot and eccentricity, and between the velocity component of W LSR and eccentricity. The eccentricities were calculated using the method described in Dinescu et al. (1999). The solar circle radius of R 0 = 8.0 kpc around the Galactic center and an LSR rotation velocity of Θ = 220 km s −1 were adopted. As expected, the total velocity is strongly correlated with eccentricity in our sample. However, the relationship between the vertical velocity and eccentricity is not obvious beyond increased scatter in V tot at large eccentricities. Several of our thick disk stars have low |W LSR | velocities ( 20 km s −1 ). Given that W LSR is proportional to z max , the maximum vertical distance of a star to the Galactic plane, these particular stars with W LSR apparently reside close to the Galactic plane.
The space motions and derived metallicities of our program stars are related. In Figure 15, we plot the three space velocity components against [Fe/H] in panels (a), (b), and (c). The larger dispersions of the U LSR , V LSR , and W LSR velocities in the true thick-disk stars is clear from inspection of these panels. Useful velocity limits in order to separate the thin and thick disk stars appear to be |W LSR | ≈ +20 km s −1 and V LSR < −40 km s −1 . However, all three velocity components that comprise total space velocity must be taken into account in assessing the Galactic membership of individual stars. In panel (d) of Figure 15 we plot [Fe/H] versus eccentricity, and it is clear that lower metallicity stars in our sample generally have more eccentric orbits. This is similar to the relation recently suggested by Lee et al. (2011). However, their study considered only thick-disk stars with subsolar metallicities ([Fe/H] < −0.3). We stress here that stars with thick-disk kinematics (including eccentricity) but [Fe/H] > −0.3 have no obvious metallicity-kinematics connections.
In Figure 16 Figure 15. They tell essentially the same story, linking in our sample the kinematics, metallicities, and α-element abundances of our sample.

EVOLUTIONARY STATUS
We estimated the probable evolutionary stages of our program stars by taking into account their loci in the T eff -log g plane, their absolute magnitudes (luminosities), and their 12 C/ 13 C values derived from our spectral analysis. In Figure 17 we show the T eff −log g diagram, using different symbols to denote stars with and without detected 13 CN spectral features. Although we selected our targets as luminosity class−III stars, our atmospheric analyses revealed about 1/3 of our sample to be either MS or SGB stars These are evident in the figure as higher temperature, higher gravity stars with weak/absent 13 CN.
In Figure 17 we have also shown theoretical evolutionary tracks from Bertelli et al. (2008) (Y=0.26, Z=0.02; note that we do have a few low metallicity stars in our sample but for simplicity we plot only tracks with a single metallicity). The evolutionary-track masses displayed here are from 0.8 to 3 M ⊙ . For tracks with 0.8−2 M ⊙ the evolutionary phases go up to RGB tip (at lower temperatures than displayed here), while the 2.5 and 3 M ⊙ tracks proceed up to the first thermal pulse stage. The base of the RGB, including masses up to 4 M ⊙ , is also shown with a dashed line, which is constructed from the data of Bertelli et al. (2008). We have also included HB evolutionary tracks for the mass range 0.55−2 M ⊙ (denoted by the thick black lines). Only 0.55 M ⊙ HB track evolves up to the early AGB phase. Other HB tracks start from the zero-age horizontal branch (ZAHB) and go up to the maximum He-burning phase. HB tracks with masses higher than 0.55 M ⊙ make an arc shape in Figure 17. This arc covers a temperature range 4500 K T eff 4900 K, and thus essentially defines the RC region. RC stars are the reddest HB stars with higher masses than normal RHB stars. Both of them are at the core He-burning stage and all belong to the RHB class in general.
In §3.1 we compared spectroscopic surface gravities with theoretical ones (Figure 4), suggesting that, for most of our targets, the masses are ∼2M ⊙ . This is consistent with the inferred evolutionary-track masses of the warmer (T eff > 5300 K), higher gravity (log g > 3.0) stars in Figure 17. The SG stars should not yet have convectively-mixed envelopes, so it is not surprising that these objects also do not have detectable 13 CN lines and hence 12 C/ 13 C > 30 (e.g. Thorén et al. 2004).
For almost all of the stars with T eff < 5300 K, the masses inferred from the evolutionary tracks in Figure 17 are 3 M ⊙ . These implied masses are much larger than those of the warmer MS/SG program stars, and also are inconsistent with our earlier ( §3.1) assertion that almost all of our stars have actual masses 2 M ⊙ . The implication is that our lower temperature stars are not SG stars on their first ascent up the RGB, but instead are post-RGB stars in the helium-burning RHB evolutionary stage. It is possible for RHBs to be confused with high-mass SGs, which can have similar absolute magnitudes to the RHBs. But such SG stars should be rare because the timescales for passage of 3 M ⊙ stars through the 4800−5300 K temperature domain (∼10 5 -10 6 years) are short compared to the RHB Heburning timescales (∼10 8 years) of less massive (thus probably more plentiful) stars. Our cooler program stars in general ought to be true members of the field RHB population.
For each program star we have estimate a probable evolutionary stage, taking into account its position in the T eff -log g plane from arguments presented above, its absolute magnitude when parallax data are available, and its 12 C/ 13 C ratio. These estimates are given in the last column of the Table 7. We suggest that there are 18 RHB stars in our sample (5 thick disk and 13 thin disk), along with several RC and RC/RHB stars, most of them members of the thin disk.

SUMMARY AND DISCUSSION
In this study we have determined atmospheric parameters and chemical compositions of field RHB candidates selected simply from their colors and absolute magnitudes. The original goal was to increase the sample size of bona-fide RHB stars, and to try to understand the physical processes involved in the existence of these relatively rare highly-evolved objects. We deliberately avoided kinematic biases by employing only photometric and spectral type information in selecting stars for observation.
We gathered high resolution, high S/N spectra of 129 candidate field RHB stars. Preliminary analyses eliminated 53 candidates from further consideration because they proved to be anomalously broad-lined rapidly rotating stars, or double-lined spectroscopic binaries, or stars for which we could not deduce reliable atmospheric parameters. For the remaining 76 program stars, we first determined values of T eff , log g, ξ t , and [Fe/H]. Then we derived abundances for α elements Si and Ca, neutron-capture elements La and Eu, proton capture elements Li, C, N, and O. We also determined 12 C/ 13 C isotopic ratios, because detection of 13 C provides strong evidence of CN-cycle H-fusion and mixing associated with evolved stars. The derived fundamental stellar parameters T eff and log g, along with the 12 C/ 13 C isotopic ratios, were used to estimate the evolutionary stages of our program stars. This evaluation suggests that about 20% of 76 program stars are true RHB members.
We then computed space velocity components for all of the program stars with available distances, proper motions, and radial velocities (either from the literature or from our own spectra). For each star we estimated Galactic population membership (thin disk, thick disk, and thin/thick transition) only from the kinematics. We examined correlations between kinematics and (a) overall metallicity, and (b) relative abundance ratios of the α elements (confirming these with abundance ratios of the neutron-capture elements). Even with our relatively small sample we recovered well-known differences between thin and thick disk stars. Our thick disk stars have higher space velocities and orbital eccentricities, lower metallicities, and larger α-element ratios than do their thin disk counterparts.
We did not take into account the kinematical constraints during the sample selection, and our moderately small survey of 76 RHB candidates turns out to contain only five probable thick disk stars. These five true thick disk RHBs are very similar to others already identified in the literature (e.g. Cottrell & Sneden 1986, Tautvaišienė et al. 2001: they are mildly metal-poor, α-enhanced (including high oxygen) stars that have low 12 C/ 13 C ratios. They also show little evidence for carbon depletion, as discussed in §4.3. Such stars have relatively low masses, and their appearance on the RHB (instead of the RC) is not surprising: mildly metal-poor globular clusters have well-populated RHBs.
A perhaps more interesting result is that our identified RHB sample is dominated by thin disk, high metallicity stars. These stars, with T eff > 5000 K, are too hot to be RC stars, which have temperatures lower than 4900 K. They also have evolved-star gravities: log g = 2.2−2.8, substantially smaller than the gravities of subgiants and main-sequence stars ( Figure 17). They show no obvious α and neutron-capture abundance anomalies compared to other thindisk samples. Their proton-capture abundances are unremarkable compared to normal thindisk RGB stars. In particular, they exhibit [O/Fe] ∼ 0 and [C/Fe] ∼ −0.4, and often low 12 C/ 13 C values, all consistent with expectations from past studies. But their residence in the RHB cannot be easily understood through standard stellar evolution considerations.
Both thick and the thin disk RHBs have a range of 12 C/ 13 C ratios ranging from 5 to 30 (Table 7). Before discussing the 12 C/ 13 C ratios of our stars in detail, we review the physical processes that can affect 12 C/ 13 C ratios during stellar evolution in the next paragraphs.
As a low mass star (< 2.25 M ⊙ ) evolves past the MS and SG evolutionary stages, the first dredge-up starts at the base of the RGB (e.g. Iben 1964Iben , 1965Iben , 1967b, accompanied by convective envelope expansion towards inner layers of the star. This will bring CN-cycle processed material up to the outer layers, thus, resulting in surface abundance alterations of the LiCNO group. Fragile Li is severely depleted first. As the convection extends into deeper inner regions, it passes through the transition region which separates the region of 12 C converted into 14 N. In standard dredge-up theories, mixing of the processed and unprocessed elements results in depletion of surface 12 C, lowering the 12 C/ 13 C ratio to ∼20−30 from its (assumed) initial solar value of ∼90 (Charbonnel 1994, Charbonnel et al. 1998, Gratton et al. 2000, while the surface 14 N values increase. First dredge−up phenomenon leaves a mean molecular gradient (µ−barrier) behind which prevents further mixing (e.g, Charbonnel et al. 1998).
However, previous observational studies have shown that 12 C/ 13 C ratios as low as the CN−cycle equilibrium value of 3.4 (e.g. Sneden et al. 1986, Cottrell & Sneden 1986, Gratton et al. 2000, Tautvaišienė et al. 2001, which obviously requires a non-canonical mixing process during the RGB phase. Recent stellar evolution studies have attempted to solve this extra-mixing problem by offering physical mechanisms such as rotation-induced mixing (Zahn 1992), cool bottom processing (CBP, Boothroyd & Sackmann 1999) and thermohaline instability plus rotation−induced mixing (Charbonnel & Zahn 2007, Charbonnel & Lagarde 2010. These ideas provide an extra-mixing process that sets in after the "µ−barrier elimination" which occurs after the convective envelope base recedes towards the surface. Then the H-burning shell eliminates the composition discontinuity left by the convective envelope and enables extra-mixing processes to come into play. For low mass stars, this stage of evolution is also called the "RGB luminosity function bump (LFB)" (Gratton et al. 2000, Charbonnel & Lagarde 2010 at which the H-burning shell burns the newly supplied fuel as it expands outwards. This stage slows down the evolution and causes a temporary luminosity drop along the RGB. For solar metallicities and initial masses of ≤ 2.25 M ⊙ LFB appears at effective temperatures lower than 4800 K (e.g. Charbonnel & Balachandran 2000, Monaco et al. 2011. As is seen in Figure 17, the evolved stars of our sample (thick+thin disk) are located approximately between 4800−5400 K in T eff and 2.2−2.8 in log g. The stars illustrated with filled circles have low 12 C/ 13 C values (≤ 20) and they clearly reside far away from LFB. We suggest that the stars with especially low isotopic ratios, 12 C/ 13 C ≤ 10, may have evolved from lower initial masses and undergone a major extra-mixing processes when they passes the LFB. The amount of the extra-mixing is related to both the initial mass and the metallicity of the star, which may explain the deviation in the 12 C/ 13 C values (≤ 20).
Stars have masses higher than 2.25 M ⊙ ignite their central He before the core becomes degenerate. The H-burning shell which surrounds the He-burning core remains until the second dredge-up and never reaches the region of molecular discontinuity left by the convective envelope during the first dredge-up phase. That is why, stars with high initial masses are not expected to show the indicators of extra-mixing, such as low 12 C/ 13 C ratios (Charbonnel 1994, Charbonnel et al. 1998, Charbonnel & Zahn 2007. In our sample, we have evolved stars with 20 < 12 C/ 13 C < 30, which are close to the values suggested by canonical models (e.g. Schaller et al. 1992, Charbonnel 1994. These stars are shown with (blue) crosses in Figure 17. Assuming that these stars did have on average larger initial masses than the stars with low isotopic ratios, then the isotopic ratio issue is "solved", but leaves a problem: why are these stars now in the RHB domain instead of the RC? Here we suggest, but cannot prove, that the RHB stars with higher 12 C/ 13 C ratios have undergone substantial mass loss at some stage(s) of their evolution, leaving them with smaller present envelope masses. Such stars will appear bluer (hotter) than RC stars.
For further insight in Figure 18 we correlate 12 C/ 13 C ratios and [Fe/H] metallicities. We include our RHB, RC, and RC/RHB stars along with the data obtained by Lambert & Ries (1981), Cottrell & Sneden (1986), Gratton et al. (2000) and Tautvaišienė et al. (2001) for the same temperature, luminosity and surface gravity range (same evolutionary stages). We have excluded SG stars and upper RGB stars in this plot. In Figure 18, we see a conspicuous trend with metallicity, similar to the one shown in Sneden (1991): lower metallicity RHB and RC stars have a smaller 12 C/ 13 C range (and lower values on average) than do similar objects of higher metallicity. As discussed in, e.g. Kalirai et al. (2009) andCatelan (2009), mass loss in evolved stars increases rapidly with increasing metallicity. If the 12 C/ 13 C ratio is also related to mass-loss along with mixing processes then for a given metallicity, stars with different mass-loss and mixing history must be responsible from the 12 C/ 13 C fluctuations seen towards higher metallicities in Figure 18.
According to Charbonnel (1994), the peak of the 13 C abundance in the inner region of a star is shifted outward towards high initial masses. If the star evolve from high-masses (> 2.25 M ⊙ ), then different mechanisms, such as winds on the RGB, for mass-loss should come into play other than He-flash. Depending on how massive the star at the beginning, it might be possible for a star to dredge up some or all of its region, where the 13 C abundance peaks, during its RGB evolution. Since the thin disk stars are more metal-rich compared to thick disk and halo stars, the thin disk stars with higher metallicities or higher masses or a combination of both might have very low 12 C/ 13 C rates or not show any 13 CN future depending on the mass fraction lost during the RGB evolution. High-mass evolving stars appear to have more options to produce a variety of surface abundances. Figure 18 should become more illuminating if we increase the number of the evolved stars which are the members of the same population. If for example we can substantially augment the RHB thick and thin disk samples, then we might have a better understanding of the mechanisms alters the CN ratios during their post-MS evolution.
We conclude by re-emphasizing our basic result: the RHB T eff -log g domain of the Galactic field is populated by sub-solar metallicity thick-disk stars (as expected) but also by high metallicity thin disk stars. We will be gathering more thin and thick disk data in order to substantially increase the RHB samples of different populations to try to set additional constraints on the mechanisms that may lead to stars ending up in this relatively rare evolutionary state.
We thank Bertrand Plez for sharing his CH line lists with us, and Andy McWilliam and Inese Ivans for making their model atmosphere interpolation software available to us. MA also would like to thank The Scientific and Technological Research Council of Turkey (TÜBİTAK) for supporting her research by the "Postdoctoral Research Scholarship" program. This research has made use of NASA's Astrophysics Data System Bibliographic Services; the SIMBAD database and the VizieR service, both operated at CDS, Strasbourg, France; and the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, CalTech, under contract with NASA. Support for this study has come from the U.S. National Science Foundation grant AST 09-08978 and the Rex G. Baker, Jr. endowment to the University of Texas Department of Astronomy.  (Feltzing & Holmberg 2000). With a V tot = 64 km s −1 , probable thick disk member with high metallicity. And probable planet host (Robinson et al. 2007).

A. Appendix
HIP 38801: [Fe/H]=-0.06. High-temperature star. Shows relatively high turbulence. It has peculiarly high La II and Eu II abundances. Has 12 C deficiency and 14 N enhancement, indicates a more evolved star than a SG but no 12 C/ 13 C ratio could be detected due to weak 13 CH or 13 CN features.