Metallicity measurements using atomic lines in M and K dwarf stars
a r X i v :a s t r o -p h /0410452v 1 19 O c t 2004
Mon.Not.R.Astron.Soc.000,1–6(2004)Printed 2February 2008
(MN L A T E X style ?le v2.2)
Metallicity measurements using atomic lines in M and K
dwarf stars
Vincent M.Woolf 1?and George Wallerstein 1?
?1
Astronomy Department,University of Washington,Box 351580,Seattle,WA 98195,USA
ABSTRACT
We report the ?rst survey of chemical abundances in M and K dwarf stars using atomic
absorption lines in high resolution spectra.We have measured Fe and Ti abundances in 35M and K dwarf stars using equivalent widths measured from λ/?λ≈33000spectra.Our analysis takes advantage of recent improvements in model atmospheres of low-temperature dwarf stars.The stars have temperatures between 3300and 4700K,with most cooler than 4100K.They cover an iron abundance range of ?2.44<[Fe /H]<+0.16.Our measurements show [Ti/Fe]decreasing with increasing [Fe/H],a trend similar to that measured for warmer stars where abundance analysis techniques have been tested more thoroughly.This study is a step toward the observational calibration of procedures to estimate the metallicity of low-mass dwarf stars using photometric and low-resolution spectral indices.
Key words:stars:abundances –stars:late-type –stars:subdwarfs.
1INTRODUCTION
Low mass (M dwarf and cooler)main sequence stars are by far the most numerous stars in the Galaxy and make up most of its baryonic mass.However,there have been few detailed chemical abundance studies of these stars with spectra of su?cient resolution for atomic absorption lines to be mea-sured individually.None of these included more than a few stars.This was largely a result of their low intrinsic luminos-ity and the molecular lines present in their spectra.Their faintness meant that few of these stars were bright enough for high signal to noise,high resolution spectra to be mea-sured easily without a large telescope.The molecular bands present in their spectra complicated the calculation of stellar model atmospheres and cause line blends which make it dif-?cult to measure atomic line strengths in large regions of the visible spectrum.There have been abundance studies which used photometry or which ?tted broad molecular features in low-resolution spectra of M and K dwarfs,but their results are not as certain as they would be if they were calibrated
?
E-mail:vmw@https://www.sodocs.net/doc/d45294209.html, (VMW);wall@https://www.sodocs.net/doc/d45294209.html, (GW)
?Based on observations obtained with the Apache Point Obser-vatory 3.5-meter telescope,which is owned and operated by the Astrophysical Research Consortium.This publication makes use of data products from the Two Micron All Sky Survey,which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technol-ogy,funded by the National Aeronautics and Space Administra-tion and the National Science Foundation.
with stars for which more precise abundances were based on higher-resolution spectra.There have also been studies which used synthetic spectrum ?tting of low resolution spec-tra to measure metallicities of low mass dwarfs.These have given rough metallicity estimates,but have not produced elemental abundances with the precision which is possible using higher resolution spectra.
Recent advances in model atmospheres of low mass dwarf stars,largely as a result of the improved treatment of molecular opacity,have provided the opportunity to de-termine abundances of M and K dwarf stars through analysis of their atomic spectral lines.Our study of chemical abun-dances in Kapteyn’s Star (HD 33793)(Woolf &Wallerstein 2004)showed us that it is possible to locate spectral regions in these stars where molecular bands are not present or are su?ciently weak for accurate equivalent widths of atomic lines to be measured.
The survey of abundance estimates we report here will provide some of the data necessary to calibrate metallicity indices for cool dwarf stars,for example the TiO and CaH indices of Reid,Hawley,&Gizis (1995)and Hawley,Gizis,&Reid (1996).With calibrated indices it should be possible to estimate the metallicity of a much larger number of cool dwarfs,including those too distant and faint to study using high resolution spectra.When a large volume-limited sample is available,it will be possible to determine if the ‘G dwarf problem’(van den Bergh 1962;Audouze &Tinsley 1976)continues to lower-mass main se-quence stars so that low-metallicity M dwarfs are more scarce than expected in our Galaxy.
2V.M.Woolf and G.Wallerstein
2OBSER V ATIONS AND REDUCTION
We selected stars for observation to cover a large range of metallicity:low metallicity stars are overrepresented in our sample compared to the actual number in the solar neighbor-hood.We increased the number of low metallicity stars by choosing to observe stars with high radial velocities,which increased the likelihood of observing halo stars,and by ob-serving stars with TiO5indices(Gizis1997)which indicate weak TiO band strengths.
Our spectrum of Kapteyn’s Star was measured with the echelle spectrograph of the4.0-m Victor M.Blanco Tele-scope at the Cerro Tololo Inter-American Observatory as described in Woolf&Wallerstein(2004).The34other stars were observed using the echelle spectrograph of the Apache Point Observatory(APO)3.5-meter telescope.
The spectra were reduced using standard IRAF routines to subtract the bias,divide by?at?eld spectra,reduce the echelle orders to one dimensional spectra,and apply ThAr lamp spectrum wavelength calibration.The spectrum of a hot,high v sin i star was used to correct for telluric absorp-tion lines where appropriate and possible.
The spectral resolution of the APO spectra is about λ/?λ≈33000,as measured using the ThAr comparison lines.There are no gaps in wavelength between echelle or-ders.The usable spectrum covers the range from about9800?A to where the measured signal from these red stars drops
o?in the blue,normally below5000?A.For most stars,the signal to noise ratio in the region of the lines we used for the analysis was at least100per pixel.It was much higher for the brighter stars(V 11)in our sample.For the faintest star,LHS364,we were only able to achieve a signal to noise ratio of about50with the observation time available.
3ANALYSIS
3.1Stellar parameters and model
As we found in our analysis of Kapteyn’s Star,the chemi-cal abundances we derive for these low-temperature dwarf stars depends strongly on the metallicity of the model at-mosphere used.For example,changing the metallicity of the model atmosphere with T e?=3500K and log g=5.0by ±0.5dex can changes the Fe abundance derived using the equivalent widths measured for Kapteyn’s Star by±0.3dex. Determining the physical parameters to use for the model atmospheres is therefore an iterative process.Fortunately, the iteration converges so it is possible to?nd the value where the metallicity derived equals the metallicity of the model atmosphere used in the analysis.
We integrated the?ux in the V,H,and K S?l-ters for a grid of synthetic spectra released with the NextGen models(Hauschildt,Allard,&Baron1999)to produce theoretical colour-temperature relation estimates. We obtained H and K S magnitudes from the2MASS point source catalog(Cutri et al.2003)and V from Mermilliod,Mermilliod,&Hauck(1997)and used these to ?nd V?K S and V?H temperatures for our stars.We used the average of the two as T e?in our models.We note that while there are few other published determinations of temperatures in the stars we observed,we were happy to see that the temperature we derive for HD88230,T e?=3970±220K,is in good agreement with the temperature derived by Ram′irez&Mel′e ndez(2004),T e?=3962±63K, using its bolometric?ux and its measured angular diameter.
For stars where parallax data was available we calcu-lated absolute H and K S magnitudes.We used these to estimate the masses using the theoretical mass-luminosity relations plotted by S′e gransan et al.(2003a).We calculated the bolometric correction BC K from the BC V and V?K NextGen colours(Hauschildt et al.1999)and used parallax, K S and BC K to derive M bol.We then used M bol,mass,and T e?to calculate log g:
log g=log M+4log(T e?/5770)+0.4(M bol?4.65)+4.44 where we have used T e?⊙=5770K,M bol⊙=4.65,log g⊙= 4.44,and M is in solar masses.For stars where parallax mea-surements were unavailable we assumed log g=5.0±0.5. Gravity does not have an e?ect on the derived chemical abundances as large as the e?ects of temperature or model metallicity:accepting this large gravity uncertainty does not produce a large uncertainty in the abundances.The paral-laxes and magnitudes used to derive the stellar physical pa-rameters are listed in Table1.Spectral types are included in the table to give a rough idea of how low resolution spectra of the stars appear:the temperatures and gravities derived for the stars and reported in Table2are more physically meaningful.
We obtained from P.Hauschildt(private communica-tion)an updated NextGen model atmosphere grid which improves on the most recent public release(Hauschildt et al. 1999)by including the improved TiO and H2O line lists de-scribed in Allard et al.(2000)in its calculation.We used this to create model atmospheres interpolated in log g and T e?for each star.
For each star we began by assuming a metallicity of [M/H]=?1.0.1The calculated Fe and Ti abundances were used to estimate the model atmosphere metallicity to be used in the next iteration.The temperature and gravity cal-culated for the stars also depend on the assumed metallic-ity,so these also varied during the iteration process.This procedure was repeated until the metallicity derived from the abundances equalled that of the model atmosphere used to calculated them.We note that in this case we de?ne the “metallicity”value by the e?ect of metals in the stellar atmo-sphere,primarily through their e?ect on continuous opacity, not by the total concentration of all elements heavier than He.
Model metallicity corrections due to non-solar[α/Fe] abundances were estimated using the LTE stellar analy-sis program MOOG(Sneden1973)output.The average Ti abundance at a star’s[Fe/H]was used as a proxy for the star’sαelement enhancement.MOOG provides partial pres-sures of requested species at the di?erent layers of the model atmosphere,so by?nding,for example,the Mg I and Mg II partial pressures we were able to estimate the fraction of Mg which is ionized at the model layer where the reference opac-ity(at1.2μm)is about0.1,approximately where the lines in which we are interested are formed.The ionization frac-tions of Na,Mg,Ca,Al,and Fe,the major electron donors, were found for each model in the grid,and thus the fraction 1We use the standard notation[X]≡log
10
(X)star?log10(X)⊙.
Metallicity measurements in M and K dwarf stars3 Table1.M and K dwarf magnitudes and parallaxes
HD33793GJ191sdM1.08.850.03 5.050.02 5.320.03255.10.9H HD36395GJ205M1.57.960.01 4.040.26 4.150.21175.7 1.2H HD88230GJ380K5 6.600.02 2.960.29 3.300.26205.20.8H HD95735GJ411M2V7.490.02 3.250.31 3.640.20392.50.9H HD97101B GJ414B M1.59.980.04 5.730.02 5.980.0283.8 1.1H HD119850GJ526M1.58.460.01 4.420.02 4.780.21184.1 1.3H HD178126G22-15K5V9.230.02 6.470.02 6.570.0241.2 1.3H HD199305GJ809M0.58.540.04 4.620.02 4.920.06142.00.8H HD217987GJ887M0.57.350.02 3.460.20 3.610.23303.90.9H LHS12HIP9560M0.512.260.048.680.028.900.0336.1 4.3H LHS38GJ412A M0.58.750.04 4.770.02 5.000.02206.9 1.2H LHS42GJ9371sdM0.012.200.038.670.028.900.0244.3 2.8H LHS104G30-48esdK713.740.0210.410.0210.570.0319.3 3.0Y LHS170HIP15234sdK10.680.017.600.027.770.0330.2 2.4H LHS173HIP16209sdK711.110.017.790.027.970.0239.2 2.5H LHS174G38-2sdM0.512.750.019.140.029.350.0222.67.4Y LHS182GJ1064D esdM0.013.900.0210.520.0210.670.0323.1 2.8Y LHS236G251-44sdK713.100.019.850.0210.000.0218.2 2.9Y LHS343G61-21sdK13.820.0210.660.0210.860.0218.6 3.7Y LHS364GJ3825esdM1.514.550.0310.860.0111.010.0236.1 3.2Y,N LHS450GJ687M39.150.03 4.550.02 4.770.03220.90.9H LHS467HIP91668esdK712.210.038.780.029.000.0226.0 3.6H LHS1138G60-18G-K13.290.0110.760.0210.920.029.5 4.6Y LHS1482L586-41K13.960.0310.820.0210.980.02
LHS1819HIP28940K410.880.028.290.038.370.0517.0 2.6H LHS1841L812-11K13.180.0310.390.0210.510.0217.5 3.3Y LHS2161G48-21K511.580.018.750.028.830.05
LHS2463G10-53K712.480.019.830.029.990.03
LHS2715GJ506.1sdK10.840.028.170.028.310.0327.9 2.5H LHS2938HIP71122K710.670.027.760.027.950.0519.0 2.0H LHS3084G15-26sdK13.430.039.780.029.990.0319.1 2.9Y LHS3356GJ701M19.370.03 5.310.02 5.570.04128.3 1.4H LHS5337G21-12M011.150.037.470.027.660.0534.5 3.3H G39-36G86-1012.360.019.540.029.760.03
HIP27928GJ9192K410.700.027.760.027.880.0326.1 2.1H
4V.M.Woolf and G.Wallerstein Table2.Fe and Ti line data
HD33793
Fe i
8047.620.86-4.65691.0 6.42 8075.150.91-5.06243.0 6.36 8327.06 2.20-1.525338.0 6.53 8387.77 2.18-1.493355.0 6.52 8514.07 2.20-2.229160.0 6.47 8515.11 3.02-2.07344.0 6.52 8582.26 2.99-2.13338.0 6.44 8611.80 2.85-1.92681.0 6.51 8621.60 2.95-2.32130.5 6.45 8674.75 2.83-1.80095.0 6.48
Metallicity measurements in M and K dwarf stars5 Table3.M and K dwarf parameters and abundances
HD337933570160 4.960.13 2.00?0.86?0.990.04?0.810.090.180.10
HD363953760140 4.710.20 1.000.150.210.130.270.130.060.08
HD882303970220 4.510.22 1.00?0.05?0.030.18?0.050.13?0.020.09
HD957353510150 4.820.24 1.00?0.40?0.420.07?0.300.090.120.08
HD97191B361040 4.650.05 1.50?0.010.020.110.080.110.060.07
HD119850365040 4.790.050.50?0.12?0.100.07?0.030.080.070.10
HD178126453030 4.570.05 1.00?0.61?0.720.07?0.340.090.380.06
HD199305372050 4.670.05 2.00?0.14?0.130.10?0.190.11?0.060.06
HD2179873680130 4.880.16 1.00?0.22?0.220.09?0.200.080.020.07
LHS12383040 4.950.140.50?0.77?0.890.04?0.720.050.170.06
LHS38360030 4.900.04 1.00?0.40?0.430.05?0.400.090.030.09
LHS42386030 5.050.090.50?0.90?1.050.04?0.890.040.160.05
LHS104397030 5.070.17 1.00?1.09?1.330.04?0.990.070.340.07
LHS170423030 4.640.10 1.00?0.81?0.970.06?0.790.050.180.04
LHS173400020 4.750.08 1.00?0.98?1.190.05?0.870.050.320.04
LHS174379020 4.780.31 1.00?0.95?1.110.05?0.830.070.280.08
LHS182387030 5.090.14 1.00?1.88?2.150.03?1.720.080.430.07
LHS236404020 4.920.16 1.50?1.07?1.320.05?1.000.060.320.06
LHS343411030 5.100.21 1.00?1.45?1.740.03?1.260.070.480.08
LHS364372030 5.440.120.50?0.82?0.930.06?0.790.070.140.08
LHS450334020 4.820.03 1.000.100.150.090.300.110.150.08
LHS467393040 4.830.15 1.25?0.93?1.100.05?0.860.070.240.06
LHS1138462030 4.860.45 1.00?2.17?2.390.04?1.930.080.460.08
LHS1482410040 5.00.5 2.00?1.59?1.880.06?1.400.090.480.12
LHS1819467040 4.540.17 1.50?0.65?0.770.09?0.250.100.520.05
LHS1841444040 5.130.20 1.50?1.18?1.470.06?1.320.070.150.05
LHS2161450030 5.00.5 1.00?0.29?0.320.08?0.130.090.190.06
LHS2463454030 5.00.5 1.50?1.62?1.890.11?1.400.090.490.10
LHS2715459040 4.850.11 1.00?0.95?1.160.05?0.710.060.450.04
LHS2938449050 4.440.13 1.50?0.20?0.210.110.000.160.210.09
LHS3084378030 4.880.16 1.00?0.64?0.730.05?0.560.050.170.05
LHS3356363030 4.790.04 1.50?0.20?0.200.08?0.250.09?0.050.05
LHS5337378040 4.590.120.50?0.45?0.500.06?0.330.050.170.06
G39-36440050 5.00.5 1.50?1.71?2.000.05?1.610.070.390.06
HIP27928437030 4.640.10 1.00?0.62?0.730.06?0.550.080.180.05
6V.M.Woolf and G.Wallerstein
to obtain low resolution spectra of the stars in our list for which TiO and CaH indices have not been measured.
By compiling a statistically signi?cant sample of ob-servationally calibrated metallicity estimates for low mass dwarfs it will be possible to determine the metallicity dis-tribution of these stars in the Galaxy.This will provide ob-servational evidence of whether there is a K and M dwarf problem in our Galaxy similar to the G dwarf problem, where low-metallicity stars are more scarce than predicted by Galactic star formation and chemical evolution models. The answer to this question,positive or negative,will pro-vide important constraints for models of the chemical en-richment of the Galaxy.
ACKNOWLEDGMENTS
We thank David Yong for help getting NextGen models to work in MOOG,Peter Hauschildt for providing pre-release NextGen atmospheres for our use,and Suzanne Hawley for helpful discussions about low mass subdwarfs.This research has made use of the SIMBAD database,operated at CDS, Strasbourg,France.This research has made use of NASA’s Astrophysics Data System Bibliographic Services.The au-thors gratefully acknowledge the?nancial support of the Kennilworth Fund of the New York Community Trust. REFERENCES
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