Chemical Abundances of 22 Extrasolar Planet Host Stars *

Mon. Not. R. Astron. Soc. 363, 71Ð78 (2005) doi:10.1111/j.1365-2966.2005.09395.x

Chemical abundances of 22 extrasolar planet host stars

C. Huang,1 G. Zhao,2† H. W. Zhang1 and Y. Q. Chen2 1Department of Astronomy, Peking University, Beijing 100871, China 2National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China

Accepted 2005 July 1. Received 2005 June 30; in original form 2005 May 19

ABSTRACT We present observations of 22 extrasolar planet host stars and derive their stellar parameters. With the high signal-to-noise ratio spectra, we acquire accurate metallicities and the differential abundances for 15 other elements and we discuss the relation between the abundance ratio and the metallicity. These sample stars are metal-rich relative to the Sun, covering the range from −0.04 to 0.54 dex with the average [Fe/H] value of 0.15 ± 0.12 dex, except for HD 37124, which has [Fe/H] =−0.45. The stars with planets show a slight overabundance pattern for [C/Fe] and [Mg/Fe], but [Na/Fe], [Al/Fe], [Si/Fe], [Ti/Fe], [Cr/Fe], [Sc/Fe], [V/Fe], [Ni/Fe] and [Ba/Fe] are approximately solar in the sample stars. These stars also show slight underabundances for [O/Fe], [Ca/Fe] and [Mn/Fe]. The sulphur displays enhanced values, ranging from −0.10 to 0.40 through the sample stars. These results are used to investigate the connection between giant planets and high metallicity and to probe the inﬂuence of the process on the other elements. Keywords: star: abundances Ð star: late-type Ð planetary systems.

et al. 1999, 2001, 2002, 2003; Gonzalez & Laws 2000; Santos, 1 INTRODUCTION Israelian & Mayor 2000; Gonzalez et al. 2001; Smith et al. 2001; In a previous investigation, Zhao et al. (2002) gave the accurate Bodaghee et al. 2003). Detailed chemical abundance analyses of the chemical abundances of 15 extrasolar planet host stars. We follow extrasolar planet host stars will provide useful information for an the method to give the ratios of [X/Fe] versus [Fe/H], which are used investigation of how the systems with large planets have formed. to search for observational connections of chemical abnormalities We compare our research work with those of other recent simi- with the presence of planets, and we investigate the nucleosynthesis lar studies to look for the trends in the data suggested in previous process of the formation of planet-harbouring systems. studies. The total number of planet-harbouring systems, which are found In this paper we present accurate metallicities of 22 extrasolar using the Doppler technique, is approaching 136 at present. The planet host stars, most of which, to some degree, were analysed in dependence of planet formation on stellar metallicity is of great previous works. The main purposes are twofold: on the one hand, we interest based on the fact that many such systems are really metal- attempt to investigate whether all our sample stars follow the planetÐ rich (see, for example, Gonzalez et al. 2001; Santos, Israelian & high-metallicity connection; on the other hand, we aim to understand Mayor 2001; Smith, Cunha & Lazzaro 2001; Sadakane, Ohkubo how the process acts on the different elements by studying the ratios & Honda 2003; Santos, Israelian & Mayor 2004) compared with between elements with the same nucleosynthesis history. the metallicity distribution of nearby field F, G, or K stars which are known to hold non-planet. However, there exist a few stars that have been found to hold planets but have quite low metal- 2 OBSERVATIONS AND DATA REDUCTIONS / =− licity, e.g. HD 114 762, which has [Fe H] 0.6. So, it seems The samples of 22 extrasolar planetary system candidates were that we cannot draw a firm conclusion and we need further stud- selected from the literature and recent radial velocity surveys of ies on other stars. Most chemical studies of planet-harbouring stars solar-type stars. The observations of 22 stars that are known to be ac- have used iron as the reference element. Some systematic stud- companied by planets were carried out by using the 2.16-m telescope ies concerning other metals, to some degree, have shown a few at the National Astronomical Observatories (Xinglong, China) with possible anomalies (see, for example, Gonzalez 1997; Sadakane the coud«e echelle spectrometer equipped with Tek CCD (1024 × 1024 pixel with 24 × 24 µ m2 each in size). Considering the low effi- ciency of the CCD in the blue region, we observed the red spectral re- Based on observations collected at the National Astronomical Observato- gion (560 <λ<900 nm) with the exception of some interorder gaps − ries (Xinglong, China). due to the limited CCD size. The 31.6 groove mm 1 R2 echelle grat- ◦ †E-mail: [email protected] ing was used and a 60 ZF3 prism was used as a cross-disperser. We

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Table 1. The stellar parameters and basic data of the sample stars.

Star VTeff log g [Fe/H] ξ t πσ(π)Mass (M)BC bÐy HD 8574 7.800 6028 4.51 0.13 1.26 22.65 0.82 1.20 −0.058 0.362 HD 9826 4.090 6156 4.14 0.15 1.39 74.25 0.72 1.34 −0.075 0.346 HD 10697 6.292 5560 3.99 0.20 1.24 30.71 0.81 1.17 −0.131 0.440 HD 16141 6.780 5642 4.09 0.07 1.47 27.85 1.39 1.09 −0.155 0.422 HD 20367 6.410 6998 4.33 0.15 1.31 36.86 1.08 1.16 −0.054 0.365 HD 23596 7.240 5885 4.09 0.31 1.41 19.24 0.85 1.20 −0.328 0.383 HD 28185 7.810 5547 4.35 0.26 1.20 25.28 1.08 0.97 −0.126 0.443 HD 33636 7.060 5875 4.49 0.03 1.14 34.85 1.33 1.10 −0.109 0.378 HD 37124 7.680 5515 4.38 −0.45 0.87 30.08 1.15 0.80 −0.150 0.421 HD 40979 6.746 6038 4.34 0.19 1.47 30.00 0.82 1.20 −0.054 0.363 HD 46375 7.840 5198 4.27 0.19 1.11 29.93 1.07 0.88 −0.333 0.502 HD 50554 6.860 5958 4.39 0.00 1.28 32.23 1.01 1.17 −0.125 0.366 HD 52265 6.301 6068 4.32 0.19 1.35 35.63 0.84 1.20 −0.059 0.360 HD 74156 7.620 5944 4.10 0.11 1.19 15.49 1.10 1.30 −0.117 0.375 HD 89744 5.741 6262 3.98 0.29 1.44 25.65 0.70 1.53 −0.199 0.338 HD 106252 7.360 5795 4.35 −0.04 1.10 26.71 0.91 1.07 −0.098 0.390 HD 168443 6.920 5455 4.00 0.10 1.16 26.40 0.85 1.07 −0.240 0.455 HD 178911 6.740 5739 3.90 −0.05 1.56 20.42 1.57 1.30 −0.180 0.403 HD 190360 5.710 5442 4.22 0.19 1.46 62.92 0.62 0.95 −0.234 0.461 HD 195019 6.910 5659 4.11 0.11 1.31 26.77 0.89 1.09 −0.158 0.419 HD 210277 6.630 5426 4.30 0.27 1.06 46.97 0.79 0.90 −0.227 0.466 HD 217107 6.180 5514 4.22 0.40 1.18 50.71 0.75 0.97 −0.298 0.456

achieved the spectral resolution around 40 000 and signal-to-noise dance analysis by requiring a zero slope of [Fe/H] versus equivalent ratio is from 150 to 250. In order to derive an accurate differential width. Finally, the whole procedure of deriving T eff, log g, [Fe/H] analysis, we also obtained a spectrum of the Moon in the red region and ξ t was iterated to consistency. The atmospheric parameters of using the same instrument set-up as that for other programme stars. 22 extrasolar planet host stars, along with the Hipparcos parallax The signal-to-noise ratio of the Moon spectrum is around 400. A and its error, mass, bolometric correction and their colour index bÐy, detail description of technical aspects of the spectrograph can be are presented in Table 1. The uncertainties of the stellar parameters found in Zhao & Li (2001). are σ (T eff) = 70 K, σ (log g) = 0.1, σ ([Fe/H]) = 0.1 and σ(ξ t) = − The data reduction was performed using the MIDAS software pack- 0.3 km s 1. age, following standard routines for order identification, background The abundance analysis is based on a net of flux-constant, ho- subtraction, order extraction, wavelength calibration, radial veloc- mogeneous, local thermodynamic equilibrium (LTE) model atmo- ity shift correction and continuum normalization. Bias, dark current spheres taken from Kurucz (1993). The abundance was calculated and scattered-light corrections are included in the background sub- with the program, ABONTEST, which was kindly supplied by Dr P. traction. The radial velocity is estimated by using a set of unblended Magain (Li«ege, Belgium). The calculations include natural broaden- lines (around 20) with intermediate strength. Then the spectrum was ing, thermal broadening, van der Waals damping and microturbulent normalized by a continuum function determined by fitting a spline Doppler broadening. ABONTEST was used to calculate the theoretical curve to a set of pre-selected continuum windows which was esti- equivalent width from the model, and elemental abundances were mated from a solar atlas. derived by requiring that the calculated equivalent widths should The equivalent widths are measured using two different methods: match those observed. Solar abundances were calculated from the direct integration of the line profile and a Gaussian function fitting. Moon spectrum, which was also observed at Xinglong Station. Usually, the latter is preferable in the case of weak lines but is Solar abundances were used to derive the stellar abundances rel- unsuitable for strong lines in which the damping wings contribute ative to the solar values. Such differential abundances are gener- significantly to the equivalent. The direct integration method gives a ally more reliable than absolute abundances because many sys- better performance for strong unblended lines. The final equivalent tematic errors nearly cancel out. Particularly, stars in this study widths are weighted averages of these two measurements, depending are similar to the Sun in their physical characteristics, and a dif- on the line intensity (see Zhao et al. 2000 for details). The typical ferential abundance analysis can avoid many possible systematic error of equivalent widths is around 4 mÅ. errors, including uncertainties in the treatment of the atmospheric structure. The log gf values we have used in this work are mostly taken 3 STELLAR PARAMETERS AND LINE from Chen et al. (2000, table 4). They start from a careful selection ANALYSIS of experiment or theoretical oscillator strengths for each element. ForCand S, which are not included in Chen et al. (2000), the The effective temperature was derived from Str¬omgren photomet- log gf value is taken from the Vienna Astrophysical Line Data ric indices (Olsen 1983, 1993) using the calibration by the infrared Base1 for C I 658.7 nm, and the gf -values of S I 604.6 and 605.2 nm flux method (Alonso, Arribas & Martinez-Roger 1996). The surface gravity was calculated from accurate Hipparcos parallaxes (ESA 1 1997). The microturbulence, ξ t ,was determined from the abun- [email protected]

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Table 2. Relative abundances with errors and the number of line used for the sample stars. For M/Fe, all abundances are given as [element/Fe], except for Fe where [Fe/H] is given.

Element [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N

HD 37124 HD 23596 HD 106252 HD 50554 HD 74156 HD 89744 C I 0.26 Ð 1 0.19 Ð 1 Ð Ð Ð 0.20 Ð 1 0.17 Ð 1 0.02 Ð 1 O I 0.15 0.05 3 Ð Ð Ð −0.29 0.05 3 −0.14 0.02 3 −0.10 0.02 3 −0.26 0.06 3 Na I 0.02 0.01 2 −0.01 0.01 2 −0.10 0.03 2 −0.13 0.02 2 −0.09 0.02 2 −0.11 0.01 2 Mg I 0.37 0.03 3 0.11 0.06 3 −0.14 0.05 2 −0.13 0.08 3 −0.08 0.04 2 0.02 0.07 2 Al I 0.13 0.08 4 −0.13 0.04 3 −0.19 0.05 4 −0.21 0.05 4 −0.02 0.05 4 −0.10 0.04 4 Si I 0.21 0.04 12 0.05 0.02 13 0.04 0.04 13 0.03 0.04 14 0.05 0.05 14 −0.02 0.06 14 S I 0.30 0.07 2 0.16 0.03 3 0.19 0.07 3 0.16 0.06 2 0.14 0.05 3 0.07 0.01 3 Ca I 0.02 0.04 12 −0.27 0.07 13 −0.21 0.04 13 −0.14 0.04 15 −0.21 0.05 12 −0.11 0.05 14 Sc II 0.12 Ð 1 −0.02 0.08 2 0.08 Ð 1 0.04 Ð 1 0.03 0.06 2 −0.12 Ð 1 Ti I 0.26 0.02 7 −0.18 0.05 7 −0.05 0.05 7 −0.02 0.02 7 −0.03 0.04 7 −0.07 0.05 6 V I 0.19 0.06 2 −0.04 0.07 2 0.07 0.05 2 0.11 0.06 2 0.06 0.08 2 0.06 0.06 2 Cr I 0.02 0.02 4 −0.05 0.05 4 −0.06 0.05 3 −0.04 0.05 3 −0.10 0.03 3 0.16 0.05 3 Mn I −0.22 0.04 2 −0.07 0.04 2 −0.21 0.03 2 −0.16 0.05 2 −0.15 0.05 2 −0.10 0.05 2 Ni I 0.00 0.04 19 −0.08 0.05 18 −0.06 0.05 18 −0.05 0.04 18 0.00 0.05 18 −0.02 0.04 19 Ba II −0.06 0.02 3 −0.13 0.03 3 0.02 0.01 4 0.03 0.06 4 0.07 0.03 4 0.17 0.05 4 Fe I −0.46 0.01 79 0.11 0.02 67 −0.13 0.02 84 −0.12 0.01 80 0.03 0.02 85 0.20 0.03 83 Fe II −0.45 0.02 8 0.31 0.05 8 −0.04 0.04 9 0.00 0.02 9 0.11 0.03 9 0.29 0.04 8

HD 28185 HD 33636 HD 52265 HD 8574 HD 9826 HD 217107 C I 0.22 Ð 1 −0.10 Ð 1 0.19 Ð 1 0.18 Ð 1 0.06 Ð 1 0.09 Ð 1 O I −0.23 0.02 3 −0.09 0.01 3 −0.11 0.01 3 −0.14 0.01 3 −0.04 0.02 3 −0.30 0.02 3 Na I 0.00 0.02 3 −0.27 0.01 3 −0.01 0.01 2 −0.21 0.05 3 −0.11 0.01 2 0.01 0.01 2 Mg I 0.13 0.04 2 −0.06 0.07 3 0.09 0.04 3 −0.05 0.03 2 0.13 0.06 3 0.07 0.06 2 Al I −0.08 0.05 4 −0.26 0.04 4 −0.08 0.05 3 −0.26 0.04 4 −0.03 0.06 3 −0.13 0.06 4 Si I 0.02 0.03 11 −0.07 0.03 12 0.16 0.05 11 0.07 0.05 10 0.08 0.04 11 0.06 0.03 12 S I 0.24 0.08 3 0.01 0.01 2 0.21 0.04 3 0.12 0.02 3 0.17 0.02 3 0.10 0.06 3 Ca I −0.22 0.05 13 −0.28 0.03 15 −0.10 0.04 12 −0.25 0.04 11 −0.13 0.05 11 −0.20 0.07 10 Sc II −0.10 Ð 1 −0.05 0.05 2 0.10 0.08 2 0.01 Ð 1 0.00 0.05 2 0.04 0.04 2 Ti I −0.11 0.07 7 −0.16 0.03 7 −0.02 0.05 7 −0.16 0.04 7 −0.03 0.04 7 −0.17 0.07 7 V I 0.04 0.08 2 0.01 0.05 2 0.15 0.06 2 0.03 0.08 2 0.08 0.05 2 0.01 0.05 2 Cr I 0.03 0.04 3 −0.11 0.03 3 0.07 0.01 3 −0.06 0.02 3 0.02 0.05 3 0.10 0.06 3 Mn I −0.08 0.05 2 −0.21 0.02 2 −0.09 0.03 2 −0.28 0.03 2 −0.12 0.04 2 −0.13 0.05 2 Ni I −0.01 0.04 18 −0.13 0.04 19 0.03 0.05 17 −0.11 0.04 19 0.00 0.04 19 −0.12 0.05 19 Ba II −0.17 0.04 4 0.14 0.03 4 0.08 0.03 3 0.15 0.06 3 0.15 0.07 3 −0.13 0.02 3 Fe I 0.11 0.03 62 −0.18 0.02 85 0.15 0.02 76 −0.07 0.02 81 0.06 0.02 78 0.20 0.05 57 Fe II 0.26 0.03 9 0.03 0.05 9 0.19 0.02 8 0.13 0.03 8 0.15 0.03 8 0.40 0.05 8

HD 210277 HD 10697 HD 16141 HD 178911 HD 20367 HD 168443 C I 0.10 Ð 1 0.26 Ð 1 0.20 Ð 1 0.46 Ð 1 −0.35 Ð 1 0.33 Ð 1 O I −0.25 0.02 3 −0.15 0.04 3 −0.01 0.03 3 0.10 0.02 3 −0.05 0.02 3 −0.04 0.02 3 Na I −0.05 0.03 2 −0.10 0.01 2 −0.08 0.01 2 0.28 0.01 2 −0.13 0.03 2 −0.06 0.03 2 Mg I 0.13 0.05 3 0.15 0.06 3 0.24 0.05 3 0.53 0.08 3 −0.04 0.05 3 0.24 0.06 3 Al I −0.12 0.04 4 −0.13 0.06 4 −0.04 0.04 4 0.20 0.06 4 −0.10 0.04 4 0.01 0.05 4 Si I −0.13 0.05 11 0.00 0.05 12 0.13 0.05 12 0.27 0.03 12 −0.01 0.03 12 0.07 0.04 13 S II 0.27 0.05 3 0.13 0.07 3 0.24 0.04 3 0.50 0.05 3 0.17 0.06 3 0.32 0.08 3 Ca I −0.24 0.05 12 −0.22 0.04 10 −0.15 0.06 13 0.11 0.03 13 −0.14 0.05 12 −0.17 0.04 11 Sc II 0.00 0.05 2 0.01 Ð 1 0.03 0.04 2 −0.03 0.05 2 −0.16 Ð 1 −0.01 Ð 1 Ti I −0.05 0.06 7 −0.19 0.05 7 −0.15 0.03 7 0.11 0.06 7 −0.14 0.03 7 −0.17 0.05 7 V I 0.09 0.06 2 −0.02 0.06 2 0.00 0.08 2 0.34 0.08 2 0.06 0.08 2 −0.04 0.08 2 Cr I 0.10 0.04 3 −0.19 0.05 4 −0.09 0.04 4 0.25 0.06 4 −0.02 0.07 4 0.05 0.04 4 Mn I −0.23 0.06 2 −0.19 0.05 2 0.04 0.03 2 0.20 0.04 2 −0.20 0.03 2 −0.26 0.04 2 Ni I −0.08 0.05 19 −0.07 0.06 19 −0.02 0.03 19 0.12 0.05 19 −0.16 0.05 19 −0.06 0.05 19 Ba II −0.06 0.05 3 0.17 0.06 3 −0.03 0.05 3 −0.03 0.06 3 0.25 0.01 3 −0.08 0.03 3 Fe I 0.15 0.03 60 0.10 0.03 65 0.01 0.03 70 0.02 0.03 70 0.01 0.02 77 −0.06 0.03 68 Fe II 0.27 0.05 8 0.20 0.04 7 0.07 0.04 7 −0.05 0.03 8 0.15 0.03 8 0.10 0.05 8 are taken from Santos et al. (2000). We emphasize that the choice of The errors of the derived abundances are estimated from the un- log gf values will not systematically affect the ﬁnal results because certainties of stellar parameters. The effects of uncertain parameters these values are adopted to derive both stellar and solar abundances, on the derived abundances are estimated by a change of 70 K in ef- and the ﬁnal stellar abundances are differential values relative to fective temperature, 0.1 dex in gravity, 0.1 dex in metallicity and the Sun. 0.3 km s−1 in microturbulence. Assuming that the individual errors

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Table 2. Ð continued

Element [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N [M/Fe] δ N

HD 195019 HD 190360 HD 40979 HD 46375 C I 0.22 ... 1 0.38 ... 1 0.11 ... 1 0.24 ... 1 O I −0.04 0.03 3 −0.04 0.02 3 0.05 0.01 3 −0.11 0.05 3 Na I −0.23 0.01 2 −0.05 0.01 2 −0.09 0.01 2 0.17 0.01 2 Mg I −0.04 0.04 2 0.17 0.08 3 −0.06 0.07 2 −0.07 0.08 2 Al I −0.13 0.07 4 0.17 0.07 4 0.01 0.05 4 0.12 0.06 4 Si I −0.01 0.04 13 0.12 0.05 13 0.12 0.04 13 0.15 0.06 12 S I 0.06 0.06 3 0.37 0.10 3 0.26 0.07 3 ...... Ca I −0.21 0.05 12 −0.19 0.04 11 −0.20 0.05 14 −0.08 0.05 13 Sc II −0.05 ... 1 0.04 0.05 2 −0.04 ... 1 0.06 ... 1 Ti I −0.21 0.04 7 −0.06 0.05 7 −0.09 0.05 7 0.05 0.05 7 V I −0.11 0.07 2 0.00 0.06 2 0.08 0.10 2 0.32 0.10 2 Cr I −0.16 0.05 4 −0.03 0.07 4 0.05 0.05 4 0.29 0.08 4 Mn I −0.32 0.05 2 −0.12 0.03 2 −0.09 0.04 2 0.03 0.06 2 Ni I −0.13 0.04 19 −0.09 0.05 19 −0.05 0.04 18 0.05 0.04 19 Ba II −0.06 0.01 3 −0.21 0.02 3 0.08 0.07 3 −0.21 0.08 3 Fe I −0.02 0.03 72 0.08 0.03 54 0.07 0.02 72 0.23 0.03 45 Fe II 0.11 0.04 8 0.19 0.03 8 0.19 0.02 7 0.19 0.04 8 are uncorrelated, we obtain the total errors of the derived abun- typical error bar of an individual element is shown in the bottom- dances, which are shown in Table 2. left corner of each panel. 16 stars of the sample show the [Fe/H] value from 0.10 to 0.40, obviously supporting the planetÐmetallicity 4 RESULTS connection theory, and ﬁve stars are similar to the Sun with the metallicity from [Fe/H] =−0.05 to 0.07. HD 37124 presents a / =− 4.1 Main results lower value with [Fe H] 0.45. Our determinations for the total of 15 elemental abundances are summarized in Table 2, which shows Fig. 1 shows the 15 different element abundances of our sample stars the relative abundances, the statistical errors of the mean abundances and the individual elemental trends with metallicity variations. The and the line numbers used for the abundance calculation.

Figure 1. Trend of [X/Fe] with [Fe/H] for C, O, Na, Mg, Al, Si, S, Ca, Sc, Ti, V, Cr, Mn, Ni and Ba. The typical error bar is shown in the bottom-left corner of each panel.

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4.2 Light elements (C, O, Na, Al) Carbon is a crucial element in the formation of ices in proto- planetary disc. It shows a overabundance pattern for [C/Fe] from the sample stars. We determine the carbon abundance from only one line, C I 658.7 nm. The results for carbon abundances should be considered to be preliminary because the line quality is not good enough for reliable abundance determination, so it is risky to draw a firm conclusion based on only one line of C I 658.7 nm. In Fig. 1, the oxygen abundances relative to the iron decrease with increasing metallicity, which is as the same as the result given by Gonzalez et al. (2001). The oxygen abundances are determined from the triplet lines 777.1, 777.4 and 777.5 nm, which are believed to suffer from non-LTE effects and are known to be very sensitive to the temperature (King 1993). We assume that the differential approach can minimize systematic error due to inappropriate temperature structures for the stars. For the solar-type stars, however, errors of the effective temperature can influence the oxygen abundances more than non-LTE effects (Takeda et al. 1998; Reetz 1999). Kiselman (2000) pointed out that the 777-nm triplet lines are not Figure 3. Plot of [Si/H] versus [Ca/H]. formed in LTE, and non-LTE effects will not change the slopes in / / the [O Fe] versus [Fe H] diagrams very much, although perhaps the α-elements keep the overabundance and tend to decrease with they will influence the error bars because of the uncertainties in the the increasing metallicity. The results are essentially the same as collisional cross-sections. Gratton et al. (1999) presented non-LTE found by Gonzalez et al. (2001). corrections to abundances of oxygen derived from LTE analyses of Magnesium and sulphur display an overabundance pattern, espe- FÐK stars over a broad range of gravities and metal abundances. Ac- cially for S with the mean value of [S/Fe] is 0.18 ± 0.12. [Mg/Fe] cording to their results, the non-LTE abundance corrections for the shows a slight decreasing trend with the increasing iron abundance O I infrared triplet are quite small ( 0.1 dex) to the solar-type dwarfs and sulphur displays a flat trend with increasing iron abundance. = = / = (Teff 6000 K, log g 4.5, [Fe H] 0.0), and LTE overestimates The behaviour of [Mg/Fe] is different from the metal-rich field the oxygen abundance. Therefore, we did not take the non-LTE star (Gonzalez et al. 2001). The [Si/Fe] values for the sample stars correction for the determination of the oxygen abundances. show a solar pattern with a plateau in the [Si/H] versus [Fe/H] As an odd-z element, Na is always thought to be primarily manu- relation, and the dispersion of the abundance ratio [Si/Fe] is very factured during supernova explosions (Timmes, Woosley & Weaver small in the metallicity range of this work. Based on the model of 1995). The abundances are based on the Na I pair near 6160 Å, and Woosley & Weaver (1995), calcium is produced in Type II super- / =− ± show a solar pattern with the mean value [ Na Fe ] 0.08 0.09. novae with intermediate mass along with silicon. Hence, calcium / [Al Fe] shows a flat trend with increasing metallicity, and most of should show the relative abundance trend as silicon in the main- the points are near the solar value. In Gonzalez et al. (2001), the sequence stars. In Fig. 3, we plot [Si/H] versus [Ca/H], which / [Al Fe] values of planet host stars are not quite separated from field shows a good relation between [Si/H] and [Ca/H] as predicted stars. theoretically. [Ti/Fe] seems to follow a continuous flat trend with increasing 4.3 α-elements (Mg, Si, S, Ca, Ti) [Fe/H] for [Fe/H] > 0.0. According to the results in Fig. 1, a clump of stars show a slight underabundance pattern. In the survey of α/ / For these elements, we plot [ Fe] versus [Fe H] in Fig. 2, where Bodaghee et al. (2003), the authors pointed out that the [Ti/Fe] α + + + is the average of Mg Si S Ca. Our results show that ratio can probably be influenced by non-LTE effects.

4.4 Scandium As an element between α-elements and iron-peak elements, scandium abundance provides some special constraints on nucleosynthesis theory. However, the observational situation for scandium is not satisfactory at present. Additionally, the significance of the hy- perfine structure (HFS) of Sc lines on the abundance pattern of Sc is not well established. Further discussion has been carried out by Nissen et al. (2000), who found that the HFS has a small influence on weak Sc II lines in metal-poor stars. Their study stops at solar metallicity, whereas we have more interest in investigating [Sc/Fe] in metal-rich stars considering HFS. In Fig. 1, we show Sc abundances considering the HFS effect; the HFS data are taken from Prochaska & McWilliam (2000). Bodaghee et al. (2003) suggest that the distribution of scandium should mimic the trend of the Figure 2. Mean α abundance [α/Fe] versus [Fe/H], where α is the average α-element Si; our results are essentially the same as in this of Mg + Si + S + Ca. discussion.

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4.5 Iron-peak elements (V, Ni, Cr, Mn) behaviour compared with Nissen et al. (2000). This result is different from the results of V and Ni; that is to say, it is due to In Fig. 1, we find that the shape of [V/Fe] resembles the function the consideration of different nucleosynthesis processes compared of silicon and scandium (Bodaghee et al. 2003). The abundances with other iron-peak elements. We calculate the Mn abundance exhibit quite a solar pattern, and overabundance at the metal-poor with the HFS effect as Sc, so it should be noted that the Mn end, and remain flat at [Fe/H] = 0.0 ∼ 0.3. Fig. 1 illustrates that the abundances derived here are somewhat lower than those of other nickel abundances show a plateau following the iron abundance with authors. The HFS data are taken from Prochaska & McWilliam small dispersion. To some degree, the [Ni/Fe] distribution exhibits (2000). avery slight underabundance (probably a solar pattern). Despite the large scatter, Fig. 1 illustrates that the [Cr/Fe] abundances remain in the solar pattern, and show a constant flat with increasing metal- 4.6 Barium licity. For the Fe-group elements, they are mainly produced by an The heavy-element abundances are very important for testing nucle- explosion of Type I supernovae. In Figs 4 and 5, we plot [V/H] ver- osynthesis theories and the model of chemical evolution of galaxies. sus [Cr/H] and [Ni/H] versus [Cr/H], which have a good one-to-one correlation. This is a strong hint of their origin, that is, they were produced from a similar nucleosynthesis process during Galactic evolution history. As an odd-z element in the Fe group, the manganese abundances increase with increasing metallicity, which show similar

Figure 4. Plot of [V/H] versus [Cr/H].

Figure 6. The comparison of stellar parameters with other works: (a1)Ð(a4) with Gonzalez et al. (2001) and Laws et al. (2003) for 14 stars; (b1)Ð(b4) Figure 5. Plot of [V/H] versus [Cr/H]. with Santos et al. (2004).

C 2005 RAS, MNRAS 363, 71Ð78 Chemical abundances of 22 extrasolar planet host stars 77

They are produced by neutron capture by two main mechanisms ac- taken from Chen et al. (2000). From the comparison of [O/Fe], cording to the strength of the neutron flux. As a neutron-capture el- [Si/Fe], [Ti/Fe], [Cr/Fe], [Ni/Fe] and [Ba/Fe] between our pro- ement, Ba is the only heavy element presented in our sample stars. gramme stars and field stars by Chen et al. (2000), the behaviours The barium distribution displays a solar pattern despite the large of the planet-harbouring stars are in close agreements with those of scatter, which has the same behaviour as normal field stars. field stars. The trend of [Na/Fe] shows a slight upward trend at high [Fe/H], which is opposite to that of Gonzalez et al. (2001) who found that the trend of [Na/Fe] in planet-harbouring stars shows a down- 5 DISCUSSION ward state at high [Fe/H]. According to our results, [Al/Fe] shows a slight underabundance compared with field stars for [Fe/H] 5.1 Comparison with previous studies 0.0. Because Al is a refractory element, if the differences of [Al/Fe] Most of our sample stars were analysed in recent spectroscopic between planet hosts and field stars are confirmed, this means that studies. There are a total 14 stars overlapping with Gonzalez and perhaps the accretion of high-Z material does not play a role in his collaborators (Gonzalez 1998a,b, 1999; Gonzalez & Laws 2000; planet host system formation. We find the signs of an overabun- Gonzalez et al. 2001; Laws et al. 2003), and total 22 stars in the new dance of [Mg/Fe] in our samples and slightly higher than field paper (Santos, Israelian & Mayor 2004). Fuhrmann (1998, 2004) stars. Zhao et al. (2002) also pointed out an overabundance of Mg made a large sample spectroscopic study for 112 thin- and thick- in their sample of 15 planet-harbouring stars. However, Sadakane disc late-type stars in the solar neighbourhood. There are three stars et al. (2002) found no sign of an overabundance of Mg in their overlapping with this study: HD 9826, HD 190360 and HD 210277. 12 samples of planet-harbouring stars. The values of [Ca/Fe] in He gave the same metallicity as that of our study within 0.05 dex. Fig. 7 show underabundance at high [Fe/H] and slightly lower than In order to have a clear picture of the present situation for these field stars. Bodaghee et al. (2003) pointed out that the [Ca/Fe] ratio extrasolar planet host stars in the studies, in Fig. 6 we plot the stellar seemed to decrease almost continuously for the entire metallicity parameters from previous quantitative analyses versus this work. It range studied in their paper, and they did not find a clearly vis- is clear that different stellar parameters and analysis methods used ible downturn in the field stars which they compared with their by different investigators lead to different results when the general sample of planet-harbouring stars. In the case of [V/Fe], we no- characteristics of the planet host stars are obtained based on the tice that planet-harbouring stars show a slight overabundance of V compiled data from different sources. to the field stars in Fig. 7. Sadakane et al. (2002) found the value of [V/Fe] for planet-harbouring stars was about 0.15 dex higher than their reference stars. Bodaghee et al. (2003) detected an over- 5.2 Comparison with field stars abundance of vanadium in metal-poor planet-harbouring stars. They In Fig. 7, we plot 11 element abundances for our 22 programme pointed out that if the differences of iron-peak element abundance stars against the metallicity compared with the data of disc stars between planet-harbouring stars and field stars are confirmed, it will

Figure 7. Trend of [X/Fe] with [Fe/H] for O, Na, Mg, Al, Si, Ca, Ti, V, Cr, Ni and Ba. Results of this work are denoted by ﬁlled triangles. Data of disc stars taken from Chen et al. (2000) are shown by plus signs.

C 2005 RAS, MNRAS 363, 71Ð78 78 C. Huang et al. support the hypothesis that planets might form more easily around Chen Y. Q., Nissen P. E., Zhao G., Zhang H. W., Benoni T., 2000, A&AS, stars with specific metallicity ratios. 141, 491 Drake J. J., 1991, MNRAS, 251, 369 ESA, 1997, The Hipparcos and Tycho Catalogues, ESA SP-1200. ESA, 6 CONCLUDING REMARKS Noordwijk Fuhrmann K., 1998, A&A, 338, 161 We obtained the chemical composition of 22 extrasolar planet host α Fuhrmann K., 2004, Astron. Nachrichten, 325, 3 stars for 15 elements including -elements (O, Mg, Si, S, Ca and Gonzalez G., 1997, MNRAS, 285, 403 Ti), Na, Sc, iron-peak elements (Cr, V and Mn) and the neutron- Gonzalez G., 1998a, A&A, 334, 221 capture element Ba. We consider the HFS effects in the studies of Gonzalez G., 1998b, A&A, 339, L29 Sc and Mn. Non-LTE effects for most light elements we have pre- Gonzalez G., 1999, ApJ, 551, L111 sented here have been investigated in some recent studies, such as Na Gonzalez G., Laws C., 2000, AJ, 119, 390 (Baum¬uller, Butler & Gehren 1998), Mg (Zhao, Butler & Gehren Gonzalez G., Laws C., Tyagi S., Reddy B. E., 2001, AJ, 121, 432 1998; Zhao & Gehren 2000), Al (Baum¬uller & Gehren 1997), S Gratton R. G., Carretta E., Eriksson K., Gustafsson B., 1999, A&A, 350, (Takada-Hidai & Takeda 1996) and Ca (Drake 1991). According 955 to these studies, non-LTE effects for the most light elements in King J. R., 1993, AJ, 106, 1206 Kiselman D., 2000, in Oxygen Abundances in Old Stars and Implications solar-type dwarf stars are likely to be insignificant. Therefore, we to Nucleosynthesis and Cosmology, 24th Meeting of the IAU, Joint have confidence that neglecting non-LTE corrections for those ele- Discussion 8, p. 5 ment abundances causes an error less than 0.05 dex. Kurucz R. L., 1993, CD-ROM 13, ATLAS9 Stellar Atmosphere Programs In connection with the abundance pattern of these elements, we and 2 kms−1 Grid. SAO, Cambridge find most elements in the sample stars are quite similar to the Sun, Laws C., Gonzalez G., Walker K., Tyagi S., Dodsworth J., Snider K., Suntzeff and there is no significant trend with metallicity separating from the N. B., 2003, AJ, 125, 2664 field stars. Because we do not have the sample stars without planets Nissen P. E., Chen Y. Q., Schuster W. J., Zhao G., 2000, A&A, 353, 722 to perform the comparison with these planet-harbouring stars, it is Olsen E. H., 1983, A&AS, 54, 55 unclear if the overabundance or underabundance trends are special Olsen E. H., 1993, A&AS, 102, 89 for the planetary systems, and future studies are desirable. Santos Prochaska J., McWilliam A., 2000, ApJ, 537, L57 Reetz J., 1999, Ap&SS, 265, 171 et al. (2001, 2004) carried out extensive spectroscopic analyses for a Sadakane K., Honda S., Kawanomoto S., Takeda Y., Takada-Hidai M., 1999, large number stars with and without planetary mass companions, and PASJ, 51, 505 obtained no significant evidence between the two types of samples. Sadakane K. et al., 2001, PASJ, 53, 315 They concluded that the metallic abundances of planet-harbouring Sadakane K., Ohkubo M., Takeda Y., Sato B., Kambe E., Aoki W., 2002, stars are essentially indistinguishable from those in usual solar-like PASJ, 54, 911 disc stars. Comparing recent results with our work, we suppose that Sadakane K., Ohkubo M., Honda S., 2003, PASJ, 55, 1005 the pollution by planet migration may not play a key role in the Santos N. C., Israelian G., Mayor M., 2000, A&A, 363, 228 observed high metallicities of planet host stars. Santos N. C., Israelian G., Mayor M., 2001, A&A, 373, 1019 Santos N. C., Israelian G., Mayor M., 2004, A&A, 415, 1153 Smith V., Cunha K., Lazzaro D., 2001, AJ, 121, 3207 ACKNOWLEDGMENTS Takada-Hidai M., Takeda Y., 1996, PASJ, 48, 739 Takeda Y., Kato K., Kawanomoto S., Sadakane K., 1998, PASJ, 50, 97 This work is supported by NKBRSF G1999075406 and the National Timmes F. X., Woosley S. E., Weaver T. A., 1995, ApJS, 98, 617 Natural Science Foundation of China under grant No. 10433010. Woosley S. E., Weaver T., 1995, ApJS, 101, 181 Zhao G., Gehren T., 2000, A&A, 362, 1077 Zhao G., Li H. B., 2001, Chin. J. Astron. Astrophys., 2, 377 REFERENCES Zhao G., Butler K., Gehren T., 1998, A&A, 333, 219 Zhao G., Qiu H. M., Chen Y. Q., Li Z. W., 2000, ApJS, 126, 461 Alonso A., Arribas S., Martinez-Roger C., 1996, A&A, 313, 873 Zhao G., Chen Y. Q., Qiu H. M., Li Z. W., 2002, AJ, 124, 2224 Baum¬uller D., Gehren T., 1997, A&A, 325, 1088 Baum¬uller D., Butler K., Gehren T., 1998, A&A, 338, 637 Bodaghee A., Santos N. C., Israelian G., Mayor M., 2003, A&A, 404, 715 This paper has been typeset from a TEX/LATEX file prepared by the author.

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