DWARFS IN THE LOCAL REGION

Transcription

1 The Astronomical Journal, 131: , 2006 June # The American Astronomical Society. All rights reserved. Printed in U.S.A. A DWARFS IN THE LOCAL REGION R. Earle Luck 1 and Ulrike Heiter 2 Received 2006 January 24; accepted 2006 March 3 ABSTRACT We present lithium, carbon, and oxygen abundance data for a sample of nearby dwarfs a total of 216 stars including samples within 15 pc of the Sun, as well as a sample of local close giant planet (CGP) hosts (55 stars) and comparison stars. The spectroscopic data for this work have a resolution of R 60; 000, a signal-to-noise ratio >150, and spectral coverage from 475 to 685 nm. We have redetermined parameters and derived additional abundances (Z > 10) for the CGP host and comparison samples. From our abundances for elements with Z > 6 we determine the mean abundance of all elements in the CGP hosts to range from 0.1 to 0.2 dex higher than nonhosts. However, when relative abundances ([x/fe]) are considered we detect no differences in the samples. We find no difference in the lithium contents of the hosts versus the nonhosts. The planet hosts appear to be the metal-rich extension of local region abundances, and overall trends in the abundances are dominated by Galactic chemical evolution. A consideration of the kinematics of the sample shows that the planet hosts are spread through velocity space; they are not exclusively stars of the thin disk. Key words: solar neighborhood stars: abundances Online material: machine-readable tables 1 Department of Astronomy, Case Western Reserve University, Euclid Avenue, Cleveland, OH ; luck@fafnir.astr.cwru.edu. 2 Department of Astronomy and Space Physics, Uppsala University, Box 515, SE Uppsala, Sweden; ulrike@astro.uu.se INTRODUCTION This paper and its predecessors (Heiter & Luck 2003; Luck & Heiter 2005, hereafter HL03 and LH05, respectively) are parts of the former NSF/ NASA Nearby Stars Project. In this context our specific aim is to examine the overall abundance properties of the local region to determine the standard of normalcy (at least in terms of abundances). We wish to do this on a very local scale, 15 pc, to examine in detail the abundance distribution. On a larger scale, out to 100 pc, we wish to statistically sample the volume using as probes both solar-type dwarfs and K giants. The primary goal is an increased understanding of the local region about the Sun. We seek the mean metallicity of the region and to determine if there are any believable temporal, spatial, or stellar characteristic related variations in the metallicity. If one considers the local region as a typical volume, then these results can be applied to other locations and thus will increase our understanding of galactic evolution. This paper further considers the 15 pc sample of LH05 and the stars of HL03 not contained in LH05. The latter stars, numbering about 100, are either close giant planet (CGP) hosts or comparison objects. It must be noted at this point that other groups are engaged in similar pursuits: notably the work presented in Allende Prieto et al. (2004), Fuhrmann (2004), and Fischer & Valenti (2005). Photometric and radial velocity studies of the local region are also in progress or have been recently completed, most notably that of Nordström et al. (2004). Given the amount of information available about these stars and the redundancy factor in the analyses, we endeavor to present our new results with brevity and point specifically at what we consider to be the crux of these new results. The parent sample for this program is those stars in the Hipparcos parallax catalog that are within 100 pc of the Sun or that are known CGP hosts (1 star in this sample is not in the Hipparcos catalog, BD , which is a CGP host). This catalog contains 22,010 stars obviously an impossibility for a highresolution spectroscopic survey. The subset of stars with an origin in the Hipparcos catalog is those stars at distances less than 15 pc from the Sun and with absolute magnitudes greater than 7.5 that are north of declination 30N0. The Hipparcos parallax catalog contains 273 stars within 15 pc. Applying the declination limit and absolute magnitude limit(s) reduces the list to about 120 stars. This sample comprises a significant fraction of all stars within 15 pc of the Sun. These stars are those that were considered in LH05. The CGP hosts are all planet hosts considered by HL03: both the ones analyzed in the work proper and those within the appendix. HL03 analyzed a comparison sample of stars drawn from lists of stars that do not have a known companion CGP (at a rather strict level) and that have spectral characteristics comparable to the CGP hosts. In addition, we gave abundances and parameters for a number of very strong lined dwarfs (that is, dwarfs imputed to be metal-rich). Here, however, we take as our two samples the totality of the available stars subdivided into CGP hosts and comparison stars. Our rationale for this lies in two statements found in Marcy et al. (2005): Among the nearest FGK main-sequence stars (d < 40 pc) the yet undiscovered giant planets typically reside beyond 1 AU because the giant planets within 1 AU have already been found and Most of the remaining undetected planets orbiting nearby FGK stars (d < 30 pc) either have masses less than 1 M Jup or reside in orbits beyond 3 AU,... Since the bulk of the comparison objects either are from the 15 pc sample or from stars known to not harbor CGPs, or fall within 40 pc of the Sun, it is unlikely that the comparison sample has any CGPs remaining within it. The possible exceptions are those stars in the mid- F-dwarf range (typically hotter than 6250 K) that are intrinsically brighter and that can reside at larger distances (up to about 80 pc). In Table 1 we present the sample of CGP hosts and related stars to be considered here and some basic information about them. Similar information about the stars of the 15 pc sample can be found in Table 1 of LH05 and here in Table 3, which gives the list of 15 pc objects along with parameters and some abundance

2 TABLE 1 Program Stars: Planet Hosts and Comparison Stars HIC HD HR Other Spectral Type m M B V Parallax d Host G0 V F8 V F H G2 IV G2 IV G5 IV H K2 V B F8 V H F H G5 V F7 V BD K5 V B G H G2 V B G H G4 IV V H G4 V H F H F8 V G0 V F H G0 III IV H G4 V G H F6 V G H ADS 6886A F8 V ADS 6886B G5 V G3/G5 V G H G3 V F9 V BDS 5037B G H BDS 5037A G F9 V G H G2 V F8 V F7 V F7 V H G H ADS 8162A K0 IV ADS 8162B K2 V H G0 V G2 III/IV G H G H K3 V B G0 V G7 V K0 V B F9 V H G5 V H F7 V H F9 V F5 V G1 V K H K0 III H A 5659 ADS 9535A G5 V B 5659 ADS 9535B G7 V G5 V H F H G0 V B G2/G3 V H F7 V

3 DWARFS IN LOCAL REGION 3071 TABLE 1 Continued HIC HD HR Other Spectral Type m M B V Parallax d Host G2 V H K0 V H G H G H F8 V H K H ADS 12101A G1 V ADS 12101B G H ADS 12169B G4 V ADS 12169A G4 V F8 V H G8 IVvar F5 V Cyg A G2 V Cyg B G5 V H G H F8 V G3/G5 III G5 IV H K H G3 V F8 V G3 IV V H ADS 14279B F7 V ADS 14279A K1 IV G6 V H F H G H F7 V ADS 16291AB F ADS 16291AB F G5 V H G8 IV H F G H BD K0 V H Notes. These are the stars from HL03 that are not within 15 pc of the Sun. The m and M columns are the apparent and absolute V magnitudes, respectively. The parallaxes are from Hipparcos ( Perryman et al ). The d column is the distance in parsecs. Host codes are ( H ) CGP host and ( B) brown dwarf host. data. All of the stars in Tables 1 and 3 have published abundances. Our goals here are (1) to analyze the remaining HL03 stars using the procedures of LH05 to homogenize the data; (2) to give detailed abundances for the stars of Table 1, spanning sodium through europium; abundances for the same elements are available for the stars of the 15 pc sample in LH05; and (3) to derive using spectrum synthesis lithium, carbon, and oxygen abundances for the complete sample. 2. SPECTROSCOPIC MATERIAL High signal-to-noise ratio (S/ N) spectra were obtained during several observing runs between 1997 and These spectra (and reductions including the equivalent width database) are the same as those used in HL03 and LH05, in which details (dates, etc.) can be found concerning the spectra. For all observations we used the Sandiford Cassegrain Echelle Spectrograph ( McCarthy et al. 1993) attached to the 2.1 m telescope at McDonald Observatory. The spectra continuously cover a wavelength range from about 484 to 700 nm, with a resolving power of about 60,000. Typical S/N values for the spectra are in excess of 150. Each night we also observed a broad-lined B star to enable cancellation of telluric lines where necessary with a S/N exceeding that of the program stars. We used IRAF 3 to perform CCD processing, scattered light subtraction, and echelle order extraction. For all further reductions a Windows-based graphical package (ASP) developed by R. E. L. was used. This includes Beer s law removal of telluric lines, smoothing with a fast Fourier transform procedure, continuum normalization, and wavelength calibration using template spectra. Finally, equivalent widths (W k ) were determined using the Gaussian approximation to the line profile. For lines with multiple measurements, average fractional differences in W k are in general lower than 15% for 10 m8 < W k < 20 m8, lower than 10% for 20 m8 < W k < 30 m8, and lower than 5% for W k > 30 m8. A comparison with published values indicates agreement at the 3% level for most stars (see Fig. 2 of HL03). Only lines with equivalent widths between 10 and 200 m8 were used for the analysis. To enable a differential analysis, we obtained a solar flux spectrum using Callisto as the reflector. We used the same spectrograph and reduction procedure as for our program stars. The measured equivalent widths are in reasonable agreement with 3 IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

4 3072 LUCK & HEITER that determined by other authors from different sources, as shown in Figure 3 of HL METHODS 3.1. Parameters and Metal Abundances This analysis follows (on the whole) the precepts laid out in HL03. The major building blocks of the analysis in addition to the spectroscopic data are atmospheric models and atomic data, starting atmospheric parameters, traditional spectroscopic criteria to determine final effective temperatures, gravities, microturbulences, and finally, the determination of elemental abundances. The model atmospheres used here were computed using the prior generation MARCS75 code (Gustafsson et al. 1975). While not a current state of the art model atmosphere generator, this code provides continuity with a wide range of previous analyses, and it has been shown that abundances generated by these models are in good agreement with those generated by other codes (HL03). As a point of reference, the model used to generate the data for the solar differential analysis was also a MARCS75 model with T ea ¼ 5770 K, log g ¼ 4:44, and V t ¼ 0:8 kms 1 (Grevesse & Sauval 1999). The atomic data for this project were assembled from a variety of laboratory sources. Bulk sources for oscillator strengths include Fuhr et al. (1988), Martin et al. (1988), and other individual sources too numerous to enumerate. These g f-values are used to determine imputed solar abundances on a per-line basis that are in turn then used to determine the differential abundances of our program stars. Line damping in F/G/K dwarfs is dominated by van der Waals damping, for which we have used the coefficients of Barklem et al. (2000) when available and have otherwise computed from the Unsöld approximation ( Unsöld 1938). Starting atmospheric parameters for the stars reanalyzed here (Table 1) were taken from our previous analysis (HL03). Final parameters were derived by enforcing traditional spectroscopic criteria for effective temperature, gravity, and microturbulence. Lines of Fe i were forced to yield zero slope in the relations between total iron abundance and excitation potential by manipulating the model effective temperature. Simultaneously, the total abundances of iron as predicated from Fe i and Fe ii were forced into equality using the model gravities. Along with the previous two forcing operations, the slope of the Fe i abundance versus equivalent width relation was minimized, with a target of zero slope. The major departure from HL03 in this study (as in LH05) is the use of an interactive line editing procedure instead of the blind statistical process used previously. The previous procedure was to use a 2 clip (computed relative to a simple mean) on the iron abundances determined from the raw data, except for the three coolest stars in the sample of HL03, for which additional lines were eliminated manually. Our current procedure is to interactively edit the raw data by visual examination of the various relations (at the best-fitting atmospheric parameters as determined from the raw data). The visual process enables one to allow for nonlinear trends in the data, determine if there are untoward trends in items such as abundance versus wavelength (indicative of a continuum problem if present), and check how elimination of lines in one relation (such as abundance vs. excitation potential) affects another (abundance vs. equivalent width). As we see below when we compare these results to those of HL03, this difference in technique leads only to minor differences in the derived abundances. Redetermined parameters and abundance data for Fe (and light-element abundances) are given in Table 2 for the CGP host stars and related objects. Table 3 carries similar information (taken from LH05 except for the light elements, which are derived here) for the 15 pc sample. The internal uncertainty in the spectroscopically determined effective temperature is of order in terms of ¼ 5040/T. This is determined from the uncertainty in the abundance versus potential slope. At 5000 K this translates to 25 K, while at 6000 K the error is 35 K. After consideration of the uncertainty due to the editing process, we feel that the actual uncertainty is more properly set at the 100 K level, and we use that as the uncertainty in all further discussions. For gravities the uncertainty is determined by consideration of how much difference can be tolerated between the total iron abundance as given by Fe i versus Fe ii. Allowing a difference no larger than 0.05 dex yields an uncertainly of 0.1 in log g. The microturbulent velocities (V t ) are determined by minimizing the slope in the Fe i abundance equivalent width relation. Nominally the minimized slope value should be 0. As shown by LH05 and Allende Prieto et al. (2004), the value of the microturbulence depends strongly on ; the sense is that as decreases, so does V t. This means that the thermal velocity grows increasingly important, finally dominating the Doppler-broadening velocity. At about 5000 K we find that the abundances of individual Fe i lines are relatively insensitive to V t : a change of 0.3 to 0.8 km s 1 in the microturbulent velocity changes the abundance by only 0.05 dex. This is fortunate, as the formal solution forcing the slope to 0 properly demands a negative microturbulence in some cases. These cases, common below T ea ¼ 5000 K, also show little change in the slope as the microturbulence varies. In HL03 we set these values to V t ¼ 0, while here we decided that 0.5 km s 1 would be more realistic. The choice fortunately has little effect on the final abundances. The final quoted abundances are differential with respect to the Sun. Mean abundances in the form [x/h], where x is a particular element (Ca, for example), are given in Tables 4 and 5 for the CGP host dwarfs and related objects. Statistics for the abundances [x/fe] and [x/h] over all stars are given in Table 6. The statistics are broken down into three groups: brown dwarf hosts, CGP hosts, and objects not known to be hosts. Details on the abundances [x/h] per species (mean, standard deviation, and number of lines) are given in Table 7. Mean and detailed abundances for the 15 pc sample can be found in LH05 Tables 5 and Syntheses To derive lithium, carbon, and oxygen abundances for our total sample, we have employed spectrum synthesis techniques. For the lithium feature we have used all components of 7 Li (using the data presented by Andersen et al. 1984) to match the observed profiles. These abundances are given in Tables 2 and 3 along with a code indicating the quality of fit: A = excellent to D = poor, while L indicates that the abundance is an upper limit. To demonstrate the quality of the fits and spectroscopic data, we show in Figure 1 three examples of spectrum fits to the lithium feature. We note that a fit to lithium has been done for our solar spectrum using a MARCS model for the Sun. Our synthesis yields a lithium abundance for the Sun of 0.85 (log, where log H ¼ 12). This is in fair agreement (considering the differences in spectra and models) with the best current value of 1.05 (Asplund et al. 2005a). To derive carbon abundances we have used C i lines at 505.2, 538.0, and nm and Swan system C 2 at nm. These lines are all of moderate strength (1 6 nm) in F/G dwarfs and thus can be synthesized with good precision. For the atomic lines we have used the oscillator strengths of Biémont et al. (1993).

5 TABLE 2 Parameter and Abundance Data for Planet Hosts and Comparison Sample ID (K) log g (cm s 2 ) V t (km s 1 ) [Fe/H] N V M (km s 1 ) Type Host Li Code h[c/h]i h[c/fe]i h[o/h]i h[o/fe]i d HR G 0.89 L HR R 1.17 L HD G H 2.54 A HR G 2.44 A HR R 0.95 L HD G H 1.35 B HD R B 0.36 L HR R H 2.00 C HD G H 2.66 A HR G 2.43 A HR R 3.15 A HIC G B 1.01 C HD G H 0.72 L HD G B 0.07 L HD G H 2.49 A HD G H 0.63 L HR G H 0.51 L HD R H 2.88 A HR R 2.61 A HR R 2.63 A HD G H 2.50 A HR G H 2.70 A HR G 1.97 A HD G H 2.20 A HR R 2.78 A HD G H 2.23 A HR G 3.22 A HR G 1.23 C HR G 1.69 A HD G H 2.46 A HR G 1.79 A HR G 2.38 A HD G H 1.07 L HD G 0.64 L HR G 2.00 A HD G H 2.32 A HR G 2.27 A HR R 1.45 L HR R 1.80 C HR R H 2.05 A HD G H 1.10 L HR G 0.60 L HD G H 0.56 L HR G 2.24 A HR G 1.17 L HD G H 1.72 B HD G H 1.00 D HD G B 0.57 D HR G 1.95 A HR G 1.40 B HD G B 0.10 L HD G H 2.03 A HR G H 1.81 A HR R H 1.91 A HR G HR R 1.77 L HR G 0.70 L HD G H 0.18 L HD G H 0.57 D HR 5659A G 0.50 L HR 5659B G 0.60 L HR G H 1.08 L HD R H 2.54 B HD R B 2.74 A

6 3074 LUCK & HEITER TABLE 2 Continued ID (K) log g (cm s 2 ) V t (km s 1 ) [Fe/H] N V M (km s 1 ) Type Host Li Code h[c/h]i h[c/fe]i h[o/h]i h[o/fe]i d HD G H 2.30 A HR R 1.00 L HR G H 1.51 C HD G H 1.20 C HD G H 0.95 C HD G H 0.58 L HR R H 0.89 L HD G H 0.64 L HR 7272A G 2.27 A HR 7272B G H 0.40 L HR G 0.48 L HR G 0.60 L HR R H 2.48 A HR G 1.00 L HD R 2.43 A HR G 1.10 D HR G H 1.20 D HD G H 0.55 L HR G 2.60 A HR G 2.22 A HD G H 1.20 A HD G H 0.12 L HR G 2.26 A HD R 0.63 L HD G H 1.44 A HR R 1.73 A HR G 0.62 C HD G H 1.50 C HD G H 2.68 A HD G H 0.80 L HR R 2.39 A HR 8687A R 2.51 B HR 8687B R 1.70 L HR G H 1.40 D HR G H 1.00 L HD R 0.07 L HD G H 0.79 L BD G H 1.02 L Notes. In these columns, V t is the microturbulent velocity, is the standard deviation about the mean of the [Fe/H] ratio, N is the number of lines (Fe i)usedinthe iron abundance determination, and V M is the macroturbulent velocity. Types are (G) Gaussian and (R) rotation profile. At V M < 5kms 1 the two profiles are indistinguishable, and we assume a Gaussian. Hosts are (H) planet host and (B) brown dwarf host. Li is the lithium abundance (or limit). The value is log Li,where log H ¼ 12. Codes are A D for the quality of fit, with A the best fit, and L denotes an abundance limit. The column h[c/ H ]i is the average [C/ H ] ratio, h[c/ Fe]iis the average carbon-to-iron ratio, is the standard deviation about the mean of the [C/H] ratio, h[o/h]i is the average [O/H] ratio, and h[o/fe]i is the average oxygen-to-iron ratio. See x 3.2 for discussion and Table 8 for individual and solar values. The d column is the dependence of the [O/H] ratio on a variation of 0.3 dex in the carbon abundance. For Swan C 2 we have used f (0; 0) ¼ 0:0303 (Grevesse et al. 1991) with the relative band f-values of Danylewych & Nicholls (1974), along with D 0 ¼ 6:210 ev (Grevesse et al. 1991) and theoretical line wavelengths (as needed) from C. Amiot (1982, private communication). We have used these data to determine solar abundances for each of these features for the Sun from our solar reflection spectrum. For these features we derive log C ¼ 8:41, 8.50, 8.54, and 8.42 (with respect to log H ¼ 12), respectively, for 505.2, 538.0, 658.7, and These abundances have been used to derive our abundances for the program stars with respect to the Sun for each feature. Note that the most recent carbon abundance determination for the Sun (Asplund et al. 2005b) yields a mean carbon abundance of They adopted the same gf-values as used here for C i but a different dissociation energy for C 2 (6.297 ev). Since we are interested primarily in differential abundances, the differences for C 2 have little to no effect on our discussion. We show in Figure 2 three fits to the C i 538 nm feature. The most remarkable thing about the data is that the C i line shows little variation in strength over the 1000 K range in effective temperature of the objects. This is actually disconcerting, as for equal abundances the lowest temperature object should show a much weaker line. In Figure 3 we show the behavior of our four carbon abundance indicators (normalized to Fe) as a function of temperature. It is obvious that C i and nm have significant problems. For the C i nm line, it is clear that the problem lies in the strong Fe i line at nm, which develops very strong wings as the temperature decreases. This makes the C i feature a marginal presence; i.e., the C i line essentially disappears into the wing of the Fe i line as the temperature decreases. For the C i 538 nm line the situation is not resolved. Since the strength of the observed line does not decrease as expected, either there is an increasing strong blend coming into play as the temperature decreases, or a strong non-lte effect comes into play. We

7 TABLE 3 Parameters and Abundance Data for the 15 pc Sample HIC HD HR (K) log g (cm s 2 ) V t (km s 1 ) [Fe/H] V M (km s 1 ) Type Host Li Code h[c/h]i h[c/fe]i h[o/h]i h[o/fe]i d G 2.28 A G H 0.59 D G 1.98 A G 0.07 A G 0.95 L R H 2.27 A G 1.86 A G 0.51 L G 0.11 L G 1.00 L G 0.79 L G 0.25 L R 2.95 A R 2.00 L G 2.67 A G 2.41 A R 2.04 A G 2.04 A G H 0.55 L G 2.39 A G 1.29 A G 0.04 L G 0.15 L R 0.79 A G 2.39 A R 2.10 L G 0.09 L G 1.98 A G 0.24 L R 2.92 A G 0.29 D G 0.07 L R 2.88 A R 2.95 A R 0.38 L G 0.25 L G 0.26 L R 1.32 L G 0.04 L G 0.93 C G 0.87 A R 2.90 A G H 0.85 L G 0.05 L R G 0.42 L R 2.53 A G 1.35 B G G H 1.75 A R 0.23 L G 0.98 B R G 0.77 B G 0.05 C G 0.63 L R 3.25 A G 1.71 A G 0.67 L R 3.10 A R 2.16 A R G 0.76 L B G 0.01 L G

8 3076 LUCK & HEITER TABLE 3 Continued HIC HD HR (K) log g (cm s 2 ) V t (km s 1 ) [Fe/H] V M (km s 1 ) Type Host Li Code h[c/h]i h[c/fe]i h[o/h]i h[o/fe]i d G 0.42 L N G 0.25 L S G 0.57 L R 1.19 L G 0.32 L G 1.08 L R 2.22 A G 0.13 L G G 0.72 C G 0.38 L G 1.92 A G 1.53 B G 1.82 A R 2.20 A G 0.28 L G 1.60 B G 0.23 L G 0.72 L G 0.39 L G 0.22 L G 0.52 L G 0.00 L G 0.12 L G 2.55 A G 0.02 L G 1.13 B G 0.08 L B G 0.14 L G 0.16 L R 2.31 A G 0.38 L G 1.29 A R 2.41 A G 0.36 L G 0.38 L G 1.00 L G 0.09 L G 0.39 L R 1.36 L G 0.98 B G 0.15 L G 0.06 L G 0.28 L R 3.08 A G 0.01 L G 0.42 L R 2.28 A G H 0.51 L G 0.50 L Notes. In these columns, V t is the microturbulent velocity, is the standard deviation about the mean of the [Fe/H] ratio, N is the number of lines (Fe i)usedintheiron abundance determination, and V M is the macroturbulent velocity. Types are (G) Gaussian and (R) rotation profile. At V M < 5kms 1 the two profiles are indistinguishable, and we assume a Gaussian. Hosts are (H) planet host and (B) brown dwarf host. Li is the lithium abundance (or limit). The value is log Li, where log H ¼ 12. Codes are A D for quality of fit, with A the best fit, and L denotes an abundance limit. The column h[c/h]i is the average [C/H] ratio, h[c/fe]i is the average carbon-to-iron ratio, is the standard deviation about the mean of the [C/H] ratio, h[o/h]i is the average [O/H] ratio, and h[o/fe]i is the average oxygen-to-iron ratio. See x 3.2fordiscussionandTable8for individual and solar values. The d column is the dependence of the [O/H] ratio on a variation of 0.3 dex in the carbon abundance. are relatively certain that the problem is not a strong non-lte effect because (1) non-lte effects are typically factors of several, not orders of magnitude as seen here; and (2) the other two C i features behave as expected, and they have excitation potentials comparable to (although they are from different levels). This leaves an intervening blend. However, we have not been able to identify the culprit. We have examined atomic line databases (NIST and VALD) without success and can find no molecular feature that is likely. To combine the carbon data to form the mean abundances given in Tables 2 and 3, we place a dividing line at an effective temperature of 5250 K. Above that temperature we use all extant data, while below that temperature we use only C 2 and C i nm.

9 TABLE 4 Average Abundances with Respect to H for Na to Co ID Na Mg Al Si S Ca Sc Ti V Cr Mn Fe Co HR HR HD HR HR HD HD HR HD HR HR BD HD HD HD HD HR HD HR HR HD HR HR HD HR HD HR HR HR HD HR HR HD HD HR HD HR HR HR HR HD HR HD HR HR HD HD HD HR HR HD HD HR HR HR HR HR HD HD HR 5659A HR 5659B HR HD HD HD