REVIEWS OF MODERN PHYSICS, VOLUME 75, JANUARY 2003

Colloquium: Stars, planets, and metals

Guillermo Gonzalez* Iowa State University, Department of Physics and Astronomy, Ames, Iowa 50011 (Published 22 January 2003)

The discovery in 1995 of the first planet orbiting another Sun-like star stimulated renewed interest in planet formation and evolution processes. A number of trends among the properties of the planets have become evident in the years since. An interesting pattern began to emerge in 1997—stars hosting planets tend to be more metal rich (i.e., have more abundant elements with ZϾ2) than the average nearby star. Other, more subtle, trends are beginning to appear as the sample size continues to grow; for example, the masses of stars hosting planets are found to correlate with their metallicities. The author reviews the state of our knowledge concerning the observed trends, their possible causes, and their possible implications for astrophysics and astrobiology.

CONTENTS B. Recommendations for future observational research 118 Acknowledgments 118 I. Introduction 101 References 118 II. Properties of Stars with Planets 102 A. The sample 102 B. Spectroscopic observations 103 C. Spectroscopic analysis methods 103 D. Photometric analysis methods 104 I. INTRODUCTION E. Results 104 1. Spectroscopic 104 The announcement of the first candidate planet orbit- 2. Photometric 104 ing another Sun-like star by Mayor and Queloz (1995) III. Comparison with Field Stars 106 opened up a new field of empirical study for astrono- A. Control samples 106 mers. The low-mass object orbiting 51 Pegasi (51 Peg), B. Comparison of metallicity distributions 107 with a minimum mass a little less than the mass of Jupi- 1. Spectroscopic 107 ter and an orbital radius of 0.05 astronomical units 2. Photometric 107 (AU’s), turned out to contradict their expectations.1 As- C. Predictions: Confirmed and not 107 tronomers had adopted the Solar System as the model D. Biases 108 for all planetary systems, having assumed that gas giant E. Cautions 109 planets would be found at least 5 AU’s from their Sun- F. Incidence of giant planets 109 like host stars. They had also expected that the planets IV. Chemical Abundance Anomalies Among Stars with Planets 110 would be found in nearly circular orbits. Both assump- A. Light elements 110 tions turned out to be wrong. 1. Lithium 110 These discoveries were made possible by great ad- 2. Beryllium 110 vances in stellar spectroscopic Doppler techniques be- B. Other elements 110 ginning in the 1980s. Instrumental in that revolution C. Trends with condensation temperature 110 were the increased use of echelle spectrographs with D. Common proper-motion pairs 111 large telescopes, which allow large-wavelength coverage V. Galactic Kinematics 111 and high spectral resolution, the development of large VI. Causes of Trends Among Stars with Planets 112 charge-couple devices (CCD’s), and the adoption of an A. Introduction 112 iodine absorption cell as the velocity reference. A giant 1. Primordial 112 planet induces typical velocity amplitude in its host star 2. Self-enrichment 112 3. Migration 114 of tens of meters per second relative to the center of B. Evaluation 114 mass of the star-planet system. Doppler velocity preci- VII. The Solar System 115 sion of 1 to 2 meters per second has been reported, but VIII. Implications and Future Research 116 3 to 5 meters per second is more typical. Achieving a A. Present research 116 Doppler precision around 1000 times better than the 1. Learning about planet formation and evolution 116

2. Distinguishing planets from brown dwarfs 117 1 3. Implications for astrophysics 117 For the sake of historical completeness, it is important to 4. Implications for astrobiology 117 note that the first substellar-mass object was found around the star HD 114762. Latham et al. (1989) reported its minimum mass as 11 times that of Jupiter and suggested it was probably a brown dwarf but possibly a giant planet. The true nature of *Electronic address: [email protected] this object is still an unsettled issue.

0034-6861/2003/75(1)/101(20)/$35.00 101 ©2003 The American Physical Society 102 Guillermo Gonzalez: Stars, planets, and metals typical stellar linewidths is possible because (1) the metallicity3 of the host star and the presence of giant wavelength coverage is large (ϳ1000 Å), and includes planets. thousands of absorption lines, (2) the signal-to-noise The first evidence purporting to link the metallicities (S/N) ratios are high (ϳ300), and variations in the point of the SWP’s with the presence of planets was published spread function of the spectrograph across a given spec- by Gonzalez (1997). Only six SWP’s with accurate me- tallicities were known at the time, but all were found to tral order are corrected by the software. have solar metallicity levels or greater; one of them, ␳1 Of course, a number of phenomena can induce appar- 55 Cnc, was known to be one of the most metal-rich ent Doppler variations in the atmosphere of a star. stars in the solar neighborhood. About the same time, These include short- and long-term variations from star Fuhrmann et al. (1997) confirmed the high metallicity of spots, pulsations, and planets. Indeed, for about two 51 Peg. The following year Gonzalez (1998a) and Fuhr- years following its announcement, a debate raged mann et al. (1998) published the results of spectroscopic around the source of the Doppler variations in 51 analyses of several more SWP’s and continued to con- Peg. Gray and Hatzes (1997) cited spectral line profile firm the high mean metallicity of the group. The close variations in their spectra of 51 Peg as evidence for non- association between the super-metal-rich class of stars radial pulsations in its atmosphere. These claims were (discussed in the literature since the 1960s) and the later disproven by Hatzes et al. (1998). This was a worth- SWP’s was also noted at the time. These types of studies while lesson. The interpretation of periodic low- have continued until the present (e.g., Gonzalez, Laws, amplitude Doppler variations around a given star as due et al., 2001; Santos et al., 2001). Although spectroscopic to the presence of a planet should be considered as pre- analyses of SWP’s have not kept pace with the rapid rate liminary until other data are obtained to eliminate alter- of planet discoveries, the sample size is now sufficient for us to begin to look for trends linking the properties nate explanations; in particlar, high-precision photomet- of the giant planets and those of their host stars. In this ric and ultraviolet Ca II emission data can be used to Colloquium, I shall review the current state of our test for pulsations and star spots, respectively. Recently, knowledge on this topic. the Doppler variations of HD 192263, previously as- cribed to the presence of a planet, were shown from II. PROPERTIES OF STARS WITH PLANETS photometry to be due to spots (Henry, Donahue, and Baliunas, 2002). A. The sample More direct methods of planet detection could even- tually move planets from the ‘‘candidate’’ column to the At the time of this writing (spring, 2002), 77 planets 4 ‘‘secure’’ column on our scorecards. One such method have been reported orbiting 69 Sun-like stars. All the would involve resolving the reflex motion of a star on discoveries to date have been accomplished with the the sky with high-precision astrometry (something that Doppler technique, whereby the reflex motion of the has yet to be achieved with any planet candidates). An- host star about the center of mass is measured over at other method would involve the detection of transits of least one orbital period of the most prominent planet. The survey stars are selected from the brightest mid-F to a planet in front of its host star via high-precision pho- M spectral type main-sequence stars5 in the solar neigh- tometry. Charbonneau et al. (2000) employed this method to confirm the planetary nature of the object in orbit around HD 209458, and the OGLE III photomet- 3 ric survey towards the Galactic bulge has recently pro- Throughout this paper the term ‘‘metallicity’’ refers to the abundances (by number) of elements heavier than He in a star duced many planet transit candidates.2 (sometimes Li is not grouped with the metals). Today, astrono- Much theoretical work has also been spurred by the mers term these elements metals because they contribute a discoveries of planets around Sun-like stars. Some of large fraction of the electrons in a stellar interior, and the met- this work has focused on the formation and dynamical als were produced only after star formation began in the uni- evolution of the planets to their present configurations verse. Historically, though, the word was chosen because the [see the review by Nelson (2001)]. Other work has fo- strong lines in most stellar spectra were known to be due to cused on trying to model the interiors and atmospheres metals, such as Ca, Na, and Fe, and astronomers of the early of substellar objects [see the review by Burrows et al. twentieth century were confident that the mix elements in the stars would be the same as in the Earth, where Fe, Ni, and Mg (2001)]. are very abundant. Unless otherwise indicated, when referring Early on, far less attention was given to another to the metallicity of a star, the relative abundances of the met- anomalous property of the 51 Peg system: the relatively als are assumed to scale with respect to the Solar System abun- high metallicity of the host star. Intrigued by this dance pattern, as determined from meteorites and the solar anomaly, the present author started a long-term project spectrum. in December 1995 to obtain high-resolution spectra of 4For the latest tabulation of extrasolar planets, see http:// stars with planets (SWP’s) as they were announced, in www.obspm.fr/encycl/encycl.html order to determine if there really is a link between the 5Here we are employing the stellar spectral type classification scheme, OBAFGKM, in order of decreasing temperature. The classification sequence was originally based on the strength of the hydrogen absorption lines in optical spectra. The Sun is 2See http://sirius.astrouw.edu.pl/˜ ogle/index.html classified as a G2 spectral type main-sequence star.

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 103 borhood, typically going out about 70 parsecs (about 230 with S/N ratios of at least several hundred per resolution light years); the samples are magnitude limited, not vol- element and resolving power near 60 000. Many SWP’s ume limited.6 To date, most SWP’s have been found by have been observed by more than one group, allowing two groups: the ‘‘California and Carnegie Planet useful cross checks of their analysis methods. Most spec- Search’’ lead by Geoffrey Marcy and Paul Butler, and tra have been obtained with 2–3-m telescopes, though the ‘‘Geneva Extrasolar Planet Search Programmes’’ led some have been obtained with slightly smaller or much by Michelle Mayor. The American group is surveying larger telescopes (e.g., the Keck 10-m telescope). about 1200 stars (Vogt et al., 2002), while the Europeans are surveying nearly 2000, with considerable overlap be- C. Spectroscopic analysis methods tween the two samples. Most of the giant planets have been found around F The methods employed to analyze the spectra of the and G dwarfs. One has been found around a giant star. SWP’s are similar to those developed over the past Three orbit subgiants. Only one M dwarf (GJ 876) is couple of decades in studies of nearby Sun-like stars. known to have planets, even though many M dwarfs are They make use of measurements of the equivalent included in the surveys. It is not clear at the present time widths (EW’s)8 of neutral and ionized iron absorption why so few planets are being found around M dwarfs. lines (Fe I, and Fe II, respectively), together with model All the target stars are sharp lined, a necessary prereq- stellar atmospheres, to derive the four basic stellar pa- uisite for the Doppler method; ␶ Boo has the broadest rameters: effective temperature (Teff), surface gravity lines, at 15 km sϪ1. ␰ (g), microturbulence velocity parameter ( t), and Due to the ambiguity of orbital inclinations, the Dop- [Fe/H].9 To date the published analyses have assumed pler method yields only minimum-mass estimates. Addi- local thermodynamic equilibrium for their stellar atmo- tional information is required to constrain the upper sphere models. mass limits. Hipparcos positional data have proven mod- Quantitative spectroscopy of Sun-like stars has im- erately useful for a few systems, but much better astrom- proved greatly over the past 40 years. Improvements in- etry is required for most (Zucker and Mazeh, 2001a). clude increased use of echelle spectrographs, better stel- Knowledge of the true mass of a given companion is lar atmosphere models, increased use of high quality needed if we are to avoid mixing into the SWP sample CCD detectors, and better measurement and analysis objects derived from fundamentally different popula- software. The uncertainty in a typical stellar [Fe/H] es- tions (e.g., brown dwarfs). timate 40 years ago was about 0.1–0.2 dex;10 this was One possible dividing line between bona fide giant reduced to about 0.05 dex by the mid 1990s. Today, the planets and brown dwarfs is the minimum mass required formal uncertainty can be as low as 0.02 dex for an old, 7 for deuterium burning, 13 MJ (Burrows et al., 2001). early G dwarf star. Strictly differential spectroscopic Even better would be a division according to formation analyses of binary companions of similar spectral types process; it could be that brown dwarfs form like stars, [e.g., Laws and Gonzalez (2001) and Gratton et al. while planets form in a protoplanetary disk. At present (2001)] can achieve formal uncertainties below 0.01 dex. we cannot be certain that there are not some brown The most critical atomic input parameter for a given dwarfs in the SWP sample. Nevertheless, the Doppler spectral line is the oscillator strength. Two approaches surveys have turned up very few objects in the mass are available with regard to its source. Either it is ob- range traditionally assumed to be occupied by brown tained from laboratory experiments or from the solar Ͻ Ͻ dwarfs (13 MJ 80). Thus some have come to call this spectrum (for which a specific set of elemental abun- the ‘‘brown dwarf desert.’’ This desert makes it less dances is adopted). To date, the spectroscopic studies likely that high-mass brown dwarfs or low-mass stars are of SWP’s have made use of solar oscillator strengths. being mistaken for giant planets. However, the situation This is the better choice, because SWP’s are similar remains ambiguous for low-mass brown dwarfs and to the Sun (in fact, the Sun is near the average of the high-mass giant planets. SWP spectral types). Use of these oscillator strengths

B. Spectroscopic observations 8The equivalent width of an absorption line is obtained by Most of the abundance analyses of SWP’s are based integrating the area under it in a spectrum wherein the con- on high-S/N-ratio, high-resolution echelle spectra. They tinuum has been normalized. The EW is the most fundamental typically cover most or all of the optical spectral range measured quantity in quantitative stellar spectroscopy. 9The bracket notation refers to the logarithmic relative num- ber abundance compared to the Sun. Thus a star with the same Fe/H abundance ratio as the Sun has ͓Fe/H͔ϭ0. Here, and 6A volume-limited sample draws from stars in a volume of very often in the literature, metallicity and [Fe/H] are em- space centered on the Sun out to some specified distance. Such ployed interchangeably, given that the abundances of the met- a sample is constructed from accurate trigonometric paral- als generally (though not always, as will be shown later) scale laxes. with Fe. 7 10 MJ is an abbreviation for the mass of Jupiter, which is about ‘‘dex,’’ the decimal exponent, is a logarithm unit that con- 0.1% the mass of the Sun (M᭪) and 318 times the mass of the verts the number before it into its common antilogarithm, e.g., 1.27 Earth (M ). 1.27 dexϭ10 .

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 104 Guillermo Gonzalez: Stars, planets, and metals minimizes several possible systematic errors resulting type. Unfortunately, spectroscopic analyses are also from deficiencies in the models or in the analysis meth- more difficult for cool dwarfs, given the increased line ods (though not completely eliminating them). Thus me- crowding in their spectra. Fortunately, there is a way tallicity determinations for SWP’s are differential with to obtain metallicity estimates for nearby K dwarfs. respect to the Sun. Kotoneva et al. (2002) have produced a metallicity Possible systematic errors include deviations from lo- calibration based on a star’s position on the cal thermodynamic equilibrium and granulation convec- Hertzsprung-Russell (H-R) diagram (a plot of luminos- tive motions, both of which have been quantified to ity versus Teff). They can obtain metallicity estimates some extent for the Sun (Holweger, 2001).11 Systematic for nearby K dwarfs to slightly better than 0.1 dex. errors are likely to be larger for SWP’s much hotter or This method works because K dwarfs burn their fuel cooler than the Sun. Highly active stellar chromo- much more slowly than more luminous stars, and thus spheres, typical of very young stars, are also likely to even the oldest ones have hardly changed their positions pose problems for quantitative spectroscopic analyses on the H-R diagram since they were born (unlike G (Gaidos and Gonzalez, 2002). These potential problems dwarfs); the tradeoff is that their ages cannot be reliably have yet to be addressed with respect to spectroscopic determined from their positions on the H-R diagram. analyses of the SWP’s. At present this is probably the best method for Due to their very crowded spectra, stars later than deriving metallicities for the few K dwarf SWP’s found mid-K spectral type have not yet been included in the to date. recent spectroscopic studies. Such stars can be analyzed in the same way as earlier spectral types, but the results will not be as precise. E. Results Once a basic set of stellar atmospheric parameters are determined for a given star, the abundances of other 1. Spectroscopic elements can be derived from individual spectral lines. Typically, two to four clean absorption lines are mea- Listed in Table I are the fundamental atmospheric pa- sured for each of 15 to 20 elements, whereas, Fe abun- rameters for SWP’s with companion masses less than dances are typically based on 40 to 60 Fe I and Fe II 13 MJ . Only the results of recent studies based on high- lines. For tabulations of abundances of elements other resolution spectroscopy are listed. Santos et al. (2001) than Fe in the SWP’s, see Gonzalez, Laws, et al. (2001) have adopted the analysis methods and linelist described and Takeda et al. (2001). by Gonzalez and Vanture (1998) and Gonzalez, Laws, et al. (2001), therefore making it more likely that the two groups yield consistent results (comparison of the results D. Photometric analysis methods for stars in common confirms this). Other studies, while employing slightly different analysis methods, yield simi- Although photometry is not a fundamental method lar results. for determining metallicities of stars (it is calibrated with Also included in the table are mass and age estimates spectroscopic estimates), it requires far less effort for a for each star. They are based on the comparison of a single star than does spectroscopic analysis. The star’s position on the H-R diagram with theoretical Stromgren uvby system is the most useful one for Sun- stellar evolutionary tracks. Useful isochrone ages can like stars. The most widely used metallicity calibration is only be derived for stars earlier than about mid-G spec- that of Schuster and Nissen (1989), which results in a tral type; the tracks converge for later spectral types. For formal uncertainty of 0.16 dex per star. Still, this calibra- the cooler stars, useful ages can be derived from Ca II tion systematically underestimates the metallicities of activity indices (which are also listed in the table). The ϳ metal-rich stars by 0.2 dex (Twarog et al., 2002). Mar- two types of age estimates generally agree when both tell and Laughlin (2002) have produced a new calibra- are available for the same stars (but there are excep- tion using a much larger sample of spectroscopic esti- tions). mates and a more complex equation; they achieve an uncertainty of 0.10 dex per star and a good fit to the metal-rich stars. 2. Photometric Photometric metallicity calibrations tend to be less reliable for dwarfs cooler than about late G spectral Gimenez (2000) published temperature and metallic- ity estimates for 25 SWP’s from uvby photometry, gen- erally confirming the spectroscopic results published un-

11 til that time. However, he did find that the photometric The derivation of absolute abundances for the solar photo- metallicities were systematically smaller by about 0.04 sphere is an important and active research program, which re- dex on average. quires the use of laboratory oscillator strengths (more on this in Sec. VII). Absolute abundances are required when, for ex- Reid (2002) also determined the scale offset (about ample, the solar photospheric abundances are to be compared 0.1 dex) between his photometrically derived metallici- with the meteoritic abundances. Absolute abundances require ties and the modern spectroscopic estimates. Such an much more careful attention to possible systematic effects than offset is likely the result of the systematic error in the do differential abundances. old uvby calibration, noted above. While such errors

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 105

TABLE I. Spectroscopic and derived properties of stars with planets.

a ␰ a Ϫ1 b Star HD Teff (K) log g t (kms ) [Fe/H] Source 1237 5520Ϯ31 4.35Ϯ0.06 1.25Ϯ0.06 ϩ0.16Ϯ0.03 6 1237 5555Ϯ50 4.65Ϯ0.15 1.50Ϯ0.08 ϩ0.11Ϯ0.08 16 4203 5587Ϯ54 4.15Ϯ0.08 1.05Ϯ0.06 ϩ0.40Ϯ0.04 10 4208 5586Ϯ34 4.39Ϯ0.04 0.69Ϯ0.10 Ϫ0.25Ϯ0.03 10 6434 5790Ϯ40 4.56Ϯ0.20 1.40Ϯ0.10 Ϫ0.55Ϯ0.07 16 8574 6034Ϯ33 4.19Ϯ0.11 1.10Ϯ0.07 ϩ0.02Ϯ0.03 10 9826 6140Ϯ60 4.12Ϯ0.11 1.35Ϯ0.10 ϩ0.12Ϯ0.05 5 9826 6107Ϯ80 4.01Ϯ0.10 1.40Ϯ0.20 ϩ0.09Ϯ0.06 3 10697 5605Ϯ36 3.96Ϯ0.07 0.95Ϯ0.07 ϩ0.16Ϯ0.03 6 12661 5714Ϯ24 4.45Ϯ0.03 0.99Ϯ0.04 ϩ0.35Ϯ0.02 6 13445 5205Ϯ30 4.70Ϯ0.07 0.82Ϯ0.06 Ϫ0.20Ϯ0.04 16 16141 5777Ϯ31 4.21Ϯ0.04 1.12Ϯ0.06 ϩ0.15Ϯ0.02 6 16141 5805Ϯ40 4.28Ϯ0.10 1.37Ϯ0.09 ϩ0.15Ϯ0.05 16 16141 5801Ϯ31 4.24Ϯ0.04 1.01Ϯ0.05 ϩ0.19Ϯ0.03 10 16141 5737Ϯ70 3.92Ϯ0.10 1.24Ϯ0.20 ϩ0.02Ϯ0.08 2 17051 6136Ϯ34 4.47Ϯ0.05 1.23Ϯ0.09 ϩ0.19Ϯ0.03 6 17051 6225Ϯ50 4.65Ϯ0.15 1.20Ϯ0.08 ϩ0.25Ϯ0.02 16 19994 6210Ϯ30 4.20Ϯ0.10 1.52Ϯ0.07 ϩ0.26Ϯ0.04 16 19994 6164Ϯ47 4.22Ϯ0.07 1.82Ϯ0.11 ϩ0.14Ϯ0.04 10 22049 5135Ϯ40 4.70Ϯ0.10 1.14Ϯ0.07 Ϫ0.07Ϯ0.06 16 22049 5086Ϯ50 4.41Ϯ0.11 0.90Ϯ0.09 Ϫ0.09Ϯ0.03 10 27442 4797Ϯ101 3.27Ϯ0.23 1.20Ϯ0.13 ϩ0.41Ϯ0.05 10 27442 4749Ϯ100 3.3Ϯ0.3 1.2Ϯ0.3 ϩ0.22Ϯ0.09 13 28185 5705Ϯ40 4.59Ϯ0.10 1.09Ϯ0.06 ϩ0.24Ϯ0.05 16 28185 5670Ϯ30 4.54Ϯ0.05 0.94Ϯ0.06 ϩ0.24Ϯ0.02 10 33636 5930Ϯ33 4.29Ϯ0.07 1.01Ϯ0.10 Ϫ0.11Ϯ0.03 10 37124 5532Ϯ39 4.56Ϯ0.04 0.85Ϯ0.11 Ϫ0.41Ϯ0.03 6 37124 5551Ϯ34 4.43Ϯ0.07 0.60Ϯ0.18 Ϫ0.37Ϯ0.03 10 38529 5646Ϯ48 3.92Ϯ0.07 1.20Ϯ0.08 ϩ0.37Ϯ0.04 6 38529 5675Ϯ40 4.01Ϯ0.15 1.39Ϯ0.09 ϩ0.39Ϯ0.06 16 38529 5500Ϯ100 3.6Ϯ0.1 1.0Ϯ0.3 ϩ0.40Ϯ0.06 15 46375 5250Ϯ55 4.44Ϯ0.08 0.80Ϯ0.07 ϩ0.21Ϯ0.04 6 46375 5241Ϯ44 4.41Ϯ0.09 0.69Ϯ0.11 ϩ0.30Ϯ0.03 10 50554 5984Ϯ31 4.37Ϯ0.05 1.04Ϯ0.09 ϩ0.02Ϯ0.02 10 52265 6162Ϯ22 4.29Ϯ0.04 1.23Ϯ0.06 ϩ0.27Ϯ0.02 6 52265 6100Ϯ30 4.29Ϯ0.10 1.31Ϯ0.05 ϩ0.24Ϯ0.05 16 75289 6140Ϯ50 4.47Ϯ0.24 1.48Ϯ0.10 ϩ0.28Ϯ0.05 5 75289 6135Ϯ40 4.43Ϯ0.20 1.50Ϯ0.07 ϩ0.27Ϯ0.06 16 75732 5250Ϯ70 4.40Ϯ0.15 0.80Ϯ0.09 ϩ0.45Ϯ0.05 7 75732 5336Ϯ90 4.47Ϯ0.10 0.76Ϯ0.20 ϩ0.40Ϯ0.07 2 80606 5645Ϯ45 4.50Ϯ0.20 0.81Ϯ0.12 ϩ0.43Ϯ0.06 12 82943 6025Ϯ40 4.54Ϯ0.10 1.10Ϯ0.07 ϩ0.33Ϯ0.06 16 82943 6008Ϯ34 4.43Ϯ0.06 1.01Ϯ0.07 ϩ0.26Ϯ0.03 10 83443 5500Ϯ60 4.50Ϯ0.20 1.12Ϯ0.09 ϩ0.39Ϯ0.09 16 83443 5389Ϯ66 4.36Ϯ0.13 0.81Ϯ0.12 ϩ0.36Ϯ0.04 10 89744 6338Ϯ39 4.17Ϯ0.05 1.55Ϯ0.09 ϩ0.30Ϯ0.03 6 92788 5775Ϯ39 4.45Ϯ0.05 1.00Ϯ0.06 ϩ0.31Ϯ0.03 6 95128 5892Ϯ70 4.27Ϯ0.10 1.01Ϯ0.20 ϩ0.00Ϯ0.07 2 95128 5861Ϯ30 4.29Ϯ0.06 1.01Ϯ0.07 ϩ0.05Ϯ0.02 10 106252 5852Ϯ31 4.39Ϯ0.05 0.97Ϯ0.09 Ϫ0.05Ϯ0.03 10 108147 6265Ϯ40 4.59Ϯ0.15 1.40Ϯ0.08 ϩ0.20Ϯ0.06 16 108147 6316Ϯ91 4.58Ϯ0.15 0.99Ϯ0.18 ϩ0.23Ϯ0.06 10 114762 5950Ϯ75 4.45Ϯ0.05 1.0Ϯ0.10 Ϫ0.60Ϯ0.06 4 117176 5500Ϯ75 3.9Ϯ0.10 1.0Ϯ0.10 Ϫ0.03Ϯ0.06 4

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TABLE I. (Continued).

a ␰ a Ϫ1 b Star HD Teff (K) log g t (kms ) [Fe/H] Source 117176 5530Ϯ45 3.99Ϯ0.05 1.05Ϯ0.05 Ϫ0.02Ϯ0.04 10 120136 6420Ϯ80 4.18Ϯ0.08 1.25Ϯ0.11 ϩ0.32Ϯ0.06 5 120136 6360Ϯ80 4.17Ϯ0.10 1.56Ϯ0.20 ϩ0.27Ϯ0.08 2 121504 6090Ϯ40 4.73Ϯ0.10 1.35Ϯ0.08 ϩ0.17Ϯ0.06 16 121504 5941Ϯ74 4.37Ϯ0.08 1.00Ϯ0.18 ϩ0.12Ϯ0.05 10 130322 5410Ϯ35 4.47Ϯ0.08 0.95Ϯ0.06 ϩ0.05Ϯ0.03 6 134987 5715Ϯ50 4.33Ϯ0.08 1.00Ϯ0.07 ϩ0.32Ϯ0.04 6 134987 5833Ϯ60 4.31Ϯ0.20 1.3Ϯ0.10 ϩ0.36Ϯ0.06 1 136118 6231Ϯ50 4.29Ϯ0.08 1.96Ϯ0.22 Ϫ0.05Ϯ0.03 10 141937 5856Ϯ49 4.44Ϯ0.08 0.96Ϯ0.06 ϩ0.14Ϯ0.04 10 143761 5821Ϯ80 4.12Ϯ0.10 1.10Ϯ0.20 Ϫ0.24Ϯ0.08 2 145675 5300Ϯ90 4.27Ϯ0.16 0.8Ϯ0.12 ϩ0.50Ϯ0.05 8 145675 5300 4.50 1.0 ϩ0.47 17 160691 5811Ϯ45 4.42Ϯ0.06 1.07Ϯ0.08 ϩ0.28Ϯ0.03 10 168443 5555Ϯ40 4.10Ϯ0.12 0.90Ϯ0.06 ϩ0.10Ϯ0.03 6 168746 5610Ϯ30 4.50Ϯ0.15 1.02Ϯ0.08 Ϫ0.06Ϯ0.05 16 168746 5577Ϯ44 4.38Ϯ0.05 0.93Ϯ0.06 Ϫ0.06Ϯ0.03 10 169830 6300Ϯ30 4.04Ϯ0.20 1.37Ϯ0.07 ϩ0.22Ϯ0.05 16 169830 6312Ϯ50 4.15Ϯ0.06 1.26Ϯ0.13 ϩ0.17Ϯ0.04 10 177830 4818Ϯ83 3.32Ϯ0.16 0.97Ϯ0.09 ϩ0.36Ϯ0.05 6 186427 5685Ϯ40 4.26Ϯ0.08 0.80Ϯ0.06 ϩ0.07Ϯ0.03 9 186427 5766Ϯ60 4.29Ϯ0.10 0.89Ϯ0.20 ϩ0.05Ϯ0.05 2 187123 5830Ϯ40 4.40Ϯ0.07 1.00Ϯ0.08 ϩ0.16Ϯ0.05 8 190228 5360Ϯ40 4.02Ϯ0.10 1.12Ϯ0.08 Ϫ0.24Ϯ0.06 16 190228 5276Ϯ6 3.51Ϯ0.06 0.99Ϯ0.05 Ϫ0.24Ϯ0.03 10 195019A 5734Ϯ32 4.09Ϯ0.10 1.10Ϯ0.05 ϩ0.03Ϯ0.03 10 209458 6063Ϯ43 4.38Ϯ0.10 1.02Ϯ0.09 ϩ0.04Ϯ0.03 6 209458 6000Ϯ50 4.25Ϯ0.20 1.15 ϩ0.00Ϯ0.02 11 210277 5540Ϯ60 4.35Ϯ0.10 0.85Ϯ0.08 ϩ0.24Ϯ0.05 8 210277 5575Ϯ30 4.44Ϯ0.10 1.12Ϯ0.08 ϩ0.23Ϯ0.05 16 210277 5541Ϯ80 4.42Ϯ0.10 0.73Ϯ0.20 ϩ0.26Ϯ0.07 2 213240 6086Ϯ83 4.51Ϯ0.16 1.00Ϯ0.13 ϩ0.23Ϯ0.06 10 217014 5795Ϯ30 4.41Ϯ0.03 1.05Ϯ0.06 ϩ0.21Ϯ0.03 6 217014 5793Ϯ70 4.33Ϯ0.10 0.95Ϯ0.20 ϩ0.20Ϯ0.07 2 217107 5600Ϯ38 4.40Ϯ0.05 0.95Ϯ0.06 ϩ0.36Ϯ0.03 6 217107 5660Ϯ30 4.42Ϯ0.05 1.01Ϯ0.06 ϩ0.39Ϯ0.04 16 217107 5597Ϯ100 4.3Ϯ0.3 0.8Ϯ0.3 ϩ0.30Ϯ0.09 13 217107 5500Ϯ100 4.25Ϯ0.10 1.0Ϯ0.3 ϩ0.31Ϯ0.10 14 222582 5735Ϯ32 4.26Ϯ0.03 0.95Ϯ0.06 ϩ0.02Ϯ0.03 6 BD-10 3166 5320Ϯ74 4.38Ϯ0.10 0.85Ϯ0.09 ϩ0.33Ϯ0.05 6 a ␰ log g is the logarithm of the surface gravity in cgs units. t is the microturbulence velocity parameter, which describes the desaturation of absorption lines by small-scale gas turbulence. bThe sources for these data are as follows: 1, Feltzing and Gustafsson, 1998; 2, Fuhrmann, 1998; 3, Fuhrmann et al., 1998; 4, Gonzalez, 1998a; 5, Gonzalez and Laws, 2000; 6, Gonzalez, Laws, et al., 2001; 7, Gonzalez and Vanture, 1998; 8, Gonzalez et al., 1999; 9, Laws and Gonzalez, 2001; 10, Laws et al., 2002; 11, Mazeh et al., 2000; 12, Naef et al., 2001; 13, Randich et al., 1999; 14, Sadakane et al., 1999; 15, Sadakane et al., 2001; 16, Santos et al., 2001; 17, Feltzing and Gonzalez, 2001.

can be corrected for, as Reid has done, their continued III. COMPARISON WITH FIELD STARS presence in modern analyses should serve as a caution to those adopting older photometric calibrations. A. Control samples As noted above, Martell and Laughlin (2002) also de- rived photometric metallicities for SWP’s, finding good The SWP sample must be placed in context in order agreement with the spectroscopic estimates. to interpret the results correctly. That context is the stel-

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 107 lar content of the same volume of space surveyed by the Perhaps the most carefully constructed field-star con- various planet-hunting groups. The larger surveys go out trol sample is that of Reid (2002). It includes 486 F, G, to about 30 parsecs, concentrating on mid-F to early-K and K stars within 25 parsecs, selected according to their dwarfs and subgiants. Therefore the ideal control Hipparcos parallaxes. sample would include all stars of the same spectral type Finally, it is important to note that some stars pres- and luminosity class in the same volume of space. Un- ently included in the control samples may eventually be fortunately, extant spectroscopic analyses fall far short found to have planets. Planets missed until now must of this goal, which is not surprising given the amount of have relatively low mass and/or long-period orbits. New work that goes into producing each quantitative stellar planet discoveries in the control samples will change the spectroscopic abundance analysis. estimates of the frequency of occurrence of SWP’s, but it The next best approach is to prepare a small sample might or might not change other trends. For example, if SWP’s, with giant planets in long-period orbits have a of nearby stars that is representative of a complete different metallicity distribution from that of the present volume-limited sample. Several such samples have been SWP’s, then present metallicity trends will be changed. employed since 1997. The quality of a given sample de- pends on several factors. Perhaps the most important one is the source of the parallax measurements used to B. Comparison of metallicity distributions select the stars. The Hipparcos parallaxes are considered 1. Spectroscopic the best overall source of parallaxes for nearby stars. Given this, samples produced after the release of the Figure 1 compares the metallicity distributions of Hipparcos data in 1997 are superior to older ones. An- SWP’s and field stars. The shape of the SWP distribution other relevant factor is the motivation for producing a is similar to that derived by Santos et al. (2001; see their given sample. Many nearby star samples have been as- Fig. 1). Several studies have applied statistical tests to sembled primarily to study Galactic chemical evolution the hypothesis that SWP and control samples were se- and not necessarily to produce a volume-limited sample. lected from the same parent populations. Most recently, Ϫ Lastly, the zero point of the metallicity scale will vary Santos et al. (2001) derived a probability of ϳ10 7 and among the various studies. Therefore any offset between Gonzalez, Laws, et al. (2001) calculated a probability of Ϫ the metallicity scales of the SWP’s and extant field-star 2.8ϫ10 6 from their respective SWP and control samples will have to be taken into account before they samples. can be compared fairly. Figure 2 compares the distributions of SWP’s and field ͓ ͔ 13 While Gonzalez (1997) did not compare the metallici- stars in Fe/H -Teff parameter space. The upper enve- ties of the few known SWP’s at that time to the field-star lope of the field-star metallicities appears to be largely distribution, it was pointed out that they overlap the independent of Teff , while that of the SWP’s increases rare class of super-metal-rich stars; this connection was towards lower Teff (though the trend is still not com- again made by Gonzalez (1998a, 1998b). In 1998b the pletely convincing). The significance of this trend will be SWP metallicities were compared to the field-star considered below. sample of Rocha-Pinto and Maciel (1996), while in 1998a and in Gonzalez, Laws, et al. (2001) they were 2. Photometric compared to stars from Favata et al. (1997) with T eff Laughlin (2000) compared his photometric metallicity Ͼ5250 K. However, Gonzalez (1998a) mistakenly em- estimates of nearby F and G dwarfs to photometric me- ployed the distribution of Favata et al. (1997) corrected tallicity estimates of the SWP’s. He found a significant for Galactic scale-height inflation. Santos, Israelian, and trend towards higher metallicity with stellar mass among Mayor (2000) employed the original Favata et al. (1997) the SWP’s relative to the field stars. Reid (2002) com- distribution, without the scale-height correction.12 pared the metallicity distribution of 68 SWP’s to that of Santos et al. (2001) prepared their own small field-star his 25-parsec control sample, confirming that the SWP’s control sample (43 stars without known planets). They are more metal rich. observed and analyzed the field stars in the same way as their SWP’s, thus ensuring the highest degree of consis- tency between the two samples. While small, this is a C. Predictions: Confirmed and not better control sample than those noted above. A control ␳1 sample size of 43 stars is adequate if all one is looking to Based on its similarity to 14 Her and 55 Cnc, do is to compare population means. Gonzalez et al. (1999) suggested that BD-10 3166 be searched for planets. Butler et al. (2000) later confirmed that a planet does indeed orbit this star. Gonzalez, Laws,

12 et al. (2001) also discuss the confirmation of their predic- The scale-height inflation corrections take into account the tion of a planet around HD 89744. Vogt et al. (2002) fact that older stars are more likely to have large velocities perpendicular to the Galactic midplane and thus will be under- represented in a volume-limited survey (since the Sun is very 13 close to the midplane). Thus, since the SWP samples are not Teff is a useful standin for the mass of a main-sequence star corrected for this bias, a field-star control sample should not be and for the mass of its outer convection zone, though the cor- corrected for it either. relation is an inverse one for the latter.

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 108 Guillermo Gonzalez: Stars, planets, and metals

FIG. 1. [Fe/H] distribution: dotted line, 43 field stars from Santos et al. (2001); solid line, 55 SWP’s from Table I; HD 4203 and BD-10 3166 are not included in the figure, since they were targeted because they were known to be metal rich. The comparable numbers in the two histograms are not indicative of the relative num- ber of stars with planets among the field stars. The two peaks in the field-star distribution are ar- tifacts of small-number statis- tics.

discovered a planet orbiting HD 4203 following the sug- Tremaine (2002) and Zucker and Mazeh (2001b) have gestion by Laughlin (2000) that it be searched because applied statistical analysis to remove this bias and derive of its high metallicity. These predictions lend a different the underlying distributions in planet mass and period kind of support for the reality of the correlation between from their observed distributions. high metallicity and the presence of giant planets. There are important biases dependent on the proper- On the other hand, Gilliland et al. (2000) failed to de- ties of the host stars. One relates to stellar age; younger tect a single planet transit in the globular cluster 47 Tuc stars exhibit faster rotation velocities and higher chro- ͓ ͔ϭϪ ( Fe/H 0.7) from their Hubble Space Telescope mospheric activity. Faster rotational velocity broadens (HST) observations. If the frequency of hot Jupiters in the spectral lines, yielding more uncertain Doppler 47 Tuc were the same as in the solar neighborhood, then shifts. In addition, higher chromospheric activity results about 17 planets should have been found in their search. in greater ‘‘velocity jitter.’’ These effects have been Low metallicity may be the dominant factor in account- quantified by the planet-hunting groups (Saar et al., ing for this null result. However, the high volume density 1998; Santos, Mayor, et al., 2000). Since the average F of stars in a globular cluster is very probably also an dwarf is younger than the average G, K, or M dwarf, this important factor, given that planet formation and sur- ‘‘youth’’ bias tends to select against the former. Yet an- vival should be sensitive to the proximity of other stars. other bias against F dwarfs is the weakness of their ab- sorption lines, due to their higher Teff values. D. Biases A bias might result from the metallicity spread among the target stars. At a given Teff , a metal-poor star will There are several important biases inherent in the have weaker spectral lines than a metal-rich one, thus Doppler planet search method. In particular, anything leading to more uncertain Doppler measurements. This that increases the uncertainty in the Doppler bias has not yet been quantified in the literature, but wavelength-shift measurements will bias the observa- Butler et al. (2000) note that ‘‘...even subsolar metallici- tions against detection of planets producing a particular ties as low as ͓Fe/H͔ϭϪ0.5 yield adequate Doppler pre- velocity amplitude. One such bias is S/N ratio, but the cision (Ϯ5msϪ1 rather than Ϯ3msϪ1) to detect nearly planet hunters are careful to adjust the integration times all planets known.’’ Most of the hot Jupiters could be to keep this quantity relatively constant. More impor- detected around stars with ͓Fe/H͔ϭϪ1. Furthermore, tantly, the Doppler method is strongly biased in favor of since the observed distribution of SWP metallicities massive planets in short-period orbits. Tabachnik and drops steeply below solar metallicity, this bias cannot

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 109

FIG. 2. The data from Table I ᭹ plotted against Teff ; , stars with planets; ϩ, field stars. Note the lack of spread of [Fe/H] among the hotter stars with planets and the larger maxi- mum [Fe/H] values among the cooler stars. See the text for discussion.

account for the preference of planets for metal-rich would like to see answered is the minimum metallicity a stars. Metal-poor stars are relatively rare in the solar star must have to be accompanied by a giant planet. This neighborhood, but this does not become an issue so long requires that the metal-poor tail of the SWP metallicity as the SWP sample is compared to a field-star control distribution be well understood. This is a difficult task, sample prepared in the same way. However, we do have because the metal-poor tail contains few SWP’s and to worry about small-number statistics for stars with hence is susceptible to contamination by other kinds of ͓ ͔ϭϪ ͓Fe/H͔рϪ0.5. objects. For example, HD 114762, with Fe/H 0.6, is Finally, bias can also result from the way the sample is sometimes included in the SWP category, but it is now selected. We have already noted that the present surveys more often classified as a brown dwarf candidate. There are biased against young and metal-poor stars. Biases may be a few other brown dwarfs scattering into the can also be introduced if the sample of space from which SWP sample, but they would have a larger systematic effect on the metal-poor part of the distribution simply a control sample is drawn differs markedly from that because the number of metal-poor SWP’s is small. included by the planet-hunting groups; the availability of Gonzalez, Laws, et al. (2001) found that two SWP’s, Hipparcos parallaxes has greatly ameliorated this pos- HD 37124 and HD 46375, are overluminous on the H-R sible source of bias. More subtle biases can result from diagram in comparison with the oldest permitted theo- color and magnitude cutoffs. Murray and Chaboyer retical evolutionary tracks. They suggested that the (2002) note that a color cutoff in a survey will be biased SWP’s might be accompanied by unresolved stellar-mass against high-mass, low-metallicity stars, given the metal- companions, which might throw off the spectroscopic licity dependence of a star’s color. A magnitude cutoff analyses. Luhman and Jayawardhana (2002) have results in a bias against low-mass, high-metallicity stars. searched for companions to these stars with adaptive They conclude that neither bias can account for the ob- optics methods and rule out companions capable of ex- served metallicity differences between SWP’s and field plaining the discrepancy. Until the source of the discrep- stars. ancy is found for these two stars, they should be treated with caution. E. Cautions F. Incidence of giant planets Because we cannot directly resolve planets orbiting other stars, we must take care with the ambiguities in Assuming observational selection biases do not domi- the dataset (as noted above). One question that many nate the SWP database, we can conclude that the aver-

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 110 Guillermo Gonzalez: Stars, planets, and metals age metallicity of the SWP distribution is higher than (2002) searched for 6Li in six field stars and eight SWP’s, that of the nearby field-star population. Tabachnik and including HD 82943, without success; however, their up- Tremaine (2002) concluded that about 3% of nearby per limit for HD 219542 A (the primary in a wide- stars have giant planets between 1 and 10 MJ and peri- separation binary system; see below) was the highest in ods between 1 and 10 years; this estimate could ap- the sample. proach 10% if we reduce the lower mass boundary to ϳ 0.5 MJ . Reid (2002) found that about 4% of the stars 2. Beryllium in his 25-parsec sample have giant planets. The inci- Several recent studies of nearby Sun-like stars have dence rises steeply with increasing metallicity, but this is shown that Li is depleted more rapidly than Be in their not to say that every metal-rich star has a giant planet. atmospheres, largely confirming theoretical expectations At ͓Fe/H͔ϭϩ0.2, only about 20% of nearby stars are that higher temperatures are needed to destroy Be accompanied by giant planets, but the incidence might (Boesgaard and King, 2002). Lopez and de Taoro (1998) approach 100% for stars with ͓Fe/H͔уϩ0.40. and Santos et al. (2002) have determined Be abundances in over 30 SWP’s, finding no significant differences be- IV. CHEMICAL ABUNDANCE ANOMALIES AMONG STARS tween SWP’s and planetless stars. WITH PLANETS

There have been several searches for abundance B. Other elements anomalies among the SWP’s. Some have claimed to have found such anomalies. We shall review the most recent Of the roughly 15 to 20 other elements included in the observational data below and consider proposed expla- published spectroscopic analyses of SWP’s, several have nations for the anomalies in a later section. been reported to display trends distinct from those ex- hibited in field stars. Gonzalez and Laws (2000) reported A. Light elements on an apparent difference in [C/Fe] among SWP’s with respect to the results of Gustafsson et al. (1999) and 1. Lithium Tomkin et al. (1997). However they later (Gonzalez, Stellar Li abundances are at once very informative Laws, et al., 2001) retracted the original claim, noting a and very difficult to interpret. This follows from the rela- systematic offset in the Gustafsson et al. (1999) results. tive delicacy of Li nuclei in the shallow interiors of F-M At the same time, Gonzalez, Laws, et al. (2001) reported dwarfs, where they are destroyed via (p,␣) reactions on differences in abundance trends for Na, Mg, and Al when they are convectively mixed into regions with compared to field stars. warm protons. Li abundances in dwarf star photo- spheres are observed to correlate most strongly with C. Trends with condensation temperature Teff , age, rotation, binarity, and metallicity. However, even within a single open cluster (such as M67), varia- The sequence of condensation of solids with distance tions in Li abundance are observed among stars that from the central star in a protoplanetary disk is depen- appear otherwise identical (Jones et al., 1999). While dent on the so-called condensation temperatures Tc of only one Li spectral feature is available for measure- the chemical elements. Values of Tc have been published ment, near 6708 Å, it is not difficult to derive reliable Li for most elements in conditions believed to have pre- abundances from high quality (i.e., high-resolution and vailed in the early solar nebula (with regard to compo- high S/N ratio) spectra of Sun-like stars. sition and pressure). This framework helps to explain Gonzalez and Laws (2000) suggested that SWP’s, the origin of the chemical differences among the con- when corrected for simple linear trends with tempera- densed bodies in the Solar System (e.g., asteroids, ter- ture, metallicity, and age, display smaller Li abundances restrial planets, gas giants). than field stars. Ryan (2000) looked at the Li abundance Strong trends with Tc have been found among the RV trends more carefully and concluded that any possible Tauri class of variable stars and attributed to grain for- differences are not significant. The larger sample size of mation and subsequent loss in a wind (see Giridhar Gonzalez, Laws, et al. (2001) confirms this conclusion. et al., 2000). Since Tc also correlates to some extent with Both 6Li and 7Li can be measured in stellar spectra, the first ionization potential of neutral atoms, it is pos- though very high quality data are needed to separate the sible that the correlation is really with the first ionization two isotopes reliably. The fully convective internal struc- potential and not with Tc . Giridhar et al. (1998) consid- ture of a typical pre-main-sequence star is expected to ered this possibility in discussing the chemical abun- completely destroy its initial allotment of 6Li. The only dance trends among the RV Tauri stars. However, they exceptions to this are metal-poor F- or early G-type found that the chemical abundances display less scatter dwarfs, for which convection zones are smaller and some when plotted against Tc . initial 6Li can be preserved (Montalban and Rebolo, Smith et al. (2001) reported that the chemical abun- 6 2002). Observational evidence for Li in some metal- dances of some SWP’s display a trend with Tc (in a sense poor stars has been reported in the literature. Israelian opposite to that seen among the RV Tauris). They made et al. (2001) reported on their detection of 6Li in the use of the abundances of 30 SWP’s reported by Gonza- atmosphere of HD 82943, a metal-rich SWP. Reddy et al. lez, Laws, et al. (2001), and references cited therein and

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FIG. 3. Plot of the slope of chemical abundances with con- densation temperature vs [Fe/ H]: ᭺, field stars; ᭹, stars with planets; FG, Feltzing and Gustafsson, 1998; BDP, Edvars- son et al., 1993 (‘‘big disk pa- per’’). The straight line is a fit to the field stars, the trend due presumably to the increase of [O/Fe] with [Fe/H] from Galac- tic chemical evolution. From Smith et al., 2001.

compared them to those of 102 field stars without known each component can be obtained. The 16 Cyg system planets from Edvardsson et al. (1993) and Feltzing and was the first common proper-motion pair in which a Gustafsson (1998). Such a comparison avoids possible planet was found (around the B component). It had systematic errors arising from the analysis methods and been known for several years that the two stars had very phenomena unrelated to the presence of planets. Their different lithium abundances, yet similar effective tem- data are reproduced in Fig. 3. The field stars display a peratures (Friel et al., 1993). Spectroscopic studies since the early 1960s had suggested that the A component was systematic increasing trend of the Tc slope with increas- ing [Fe/H]. Smith et al. attribute the bulk of the field-star slightly more metal rich than the B component, but the trend to the well-established trend of [O/Fe] with [Fe/H] uncertainties were too large to show this conclusively. due to Galactic chemical evolution; the O abundance Laws and Gonzalez (2001), applying a differential abun- dance analysis method, found a significant difference in carries much weight in this type of analysis, given that [Fe/H] between the two stars; the A component is more few elements have comparably small values of T . They c iron rich than the B component by 0.025Ϯ0.009 dex. found that five or six SWP’s deviate from the field-star Their relative abundances also display a weak trend with trend (and from the other SWP’s) on the positive side. Tc . Gratton et al. (2001) performed differential abun- The magnitude of the effect is not large, so their result dance analyses, patterned after those of Laws and should be considered as very tentative. It is also inter- Gonzalez (2001) on six common proper-motion pairs not esting that this small group consists mostly of hot stars known to have planets and found a significant trend with with close-in planets. There is also evidence of trends Tc in the HD 219542 A and B pair. with Tc among a couple of common proper-motion pairs Zucker et al. (2002) have reported on the planet can- (see below). didate in orbit about HD 178911 B, which is part of a Takeda et al. (2001) also searched for a correlation triple star system. The brighter companion, HD 178911 between chemical abundances and Tc . They performed A, is itself a close binary. HD 178911 A and B are well independent analyses on 14 SWP’s and 4 planetless field separated in the sky, allowing spectra uncontaminated stars, finding no significant trends with Tc . Unfortu- from the companion to be obtained for each star. The nately, many more field stars than Takeda et al. (2001) two components of the primary are separated by only observed would be required to perform a convincing about 3 AU’s, which is unresolvable. Like the 16 Cyg test. Thus, at this time, the chemical abundance trends system, HD 178911 B has a much smaller Li abundance among the chemical elements in single SWP’s are sug- than the A component. However, this system differs gestive, but more data are required. from the 16 Cyg one in that the components of HD 178911 A are too close together to allow the formation of a giant planet in a moderate-sized orbit. This has im- D. Common proper-motion pairs plications as to the possible explanation for the different lithium abundances. Common proper-motion pairs are wide-separation bi- nary stars, usually discovered in astrometric proper- V. GALACTIC KINEMATICS motion surveys. Stars in common proper-motion pairs presumably formed together out of the same birth cloud. Galactic kinematics is a useful observational discrimi- Many are sufficiently far apart so that clean spectra of nant for various populations of objects in the Milky Way

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 112 Guillermo Gonzalez: Stars, planets, and metals

Galaxy. Older, metal-poor stars tend to have ‘‘hotter’’ method. Since the first planets to be discovered in a kinematics. Gonzalez (1999) compared the kinematics of given survey are those with the shortest orbital periods, 11 SWP’s to field stars (mostly from Edvardsson et al., then the high mean metallicity of the SWP sample 1993) and found that the SWP’s tend to be more metal should decrease as the time baseline increases and un- rich, even when corrected for differences in ages and covers planets in larger orbits that had not migrated. We Galactic positions. The youngest star in the sample, ␶ shall call this second explanation migration. Boo was found to be more metal rich than any young Gonzalez (1998a, 1998b) added a third possibility; gi- star in the Edvardsson et al. (1993) sample. Gonzalez, ant planets might be more likely to form around a star Laws, et al. (2001) revisited the kinematics of the young that forms from a metal-rich interstellar natal cloud. SWP’s with a larger sample and found that the SWP’s This would be expected if giant planets form from an tend to have higher metallicities than field stars at the initial rock/ice core onto which abundant hydrogen and same mean Galactocentric distance. helium gas accrete (Podolak et al., 1993). The process of Most recently, Barbieri and Gratton (2002) have com- core formation is expected to operate more rapidly pared the kinematics of 58 SWP’s with the Edvardsson when the surface density of solid material is initially et al. (1993) sample. They found that the mean metallic- greater (Pollack et al., 1996). We shall call this explana- ity of the SWP’s is greater than that of the field stars at tion primordial. Of course, these three classes of expla- every perigalactic distance. nations are not mutually exclusive. The evidence cited for each of these three explanations is summarized be- low and some evaluations are offered afterwards. VI. CAUSES OF TRENDS AMONG STARS WITH PLANETS

A. Introduction 1. Primordial

Why does the incidence of giant planets rise so steeply The evidence cited in favor of the primordial explana- towards higher metallicity? Gonzalez (1997) proposed tion includes some very metal-rich SWP’s with deep con- two possible explanations for this trend. First, in the pro- vection zones. The most extreme examples are ␳1 55 cess of planet formation and evolution, some circumstel- Cnc and 14 Her, which are on the lower main sequence lar matter enriched in metals might mix with the outer (and thus have thick outer convection zones) and are the convective envelope of the host star. The accreted mat- most metal-rich stars known in the solar neighborhood. ter could come in several forms: terrestrial planets, as- In addition, HD 27442, HD 38529, and HD 177830 are teroids and comets, and gas giants. To have an observ- very metal-rich subgiants. A subgiant experiences a able effect, the accretion must occur after the star rapid increase in the depth of its outer convection zone reaches the main sequence and its outer convection zone as it leaves the main sequence. Any early enrichment of shrinks to a small fraction of its total mass. The accreted its convective envelope would be greatly diluted when it matter should mix only with the convective region of the becomes a subgiant. While these cases rule out self- star. Thus, within this scenario, a star accompanied by enrichment as the cause of their high metallicities, the planets is more likely to have the metallicity of its con- migration explanation cannot be ruled out. vective envelope increased above its initial value. We shall call this explanation self-enrichment.14 2. Self-enrichment The second explanation is based on the mechanism of planet migration, which was originally proposed to ac- By the time the Sun was a few hundred million years count for the very short orbital period of the giant old, its outer convection zone had shrunk to about planet around 51 Peg (Lin et al., 1996). Such a planet 0.03 M᭪ . For a 1.2-M᭪ star, the outer convection zone is presumably migrated from at least a few AU’s away (be- only 0.006 M᭪ at the same age. Thus the sensitivity of yond the water-ice condensation boundary) to its the surface metallicity to accreted matter rises steeply present location near 0.05 AU. Migration might result for stars only slightly more massive than the Sun. The from gravitational interactions with gas or condensed addition of half an Earth mass of iron to the Sun would bodies in the disk, or it might occur via mutual gravita- have increased its [Fe/H] by about 0.017 dex (Murray tional encounters between two or more giant planets et al., 2001); see also Fig. 5 of Ford et al. (1999) for esti- (Rasio and Ford, 1996; Weidenschilling and Marzari, mates of the increase of [Fe/H] in the solar atmosphere 1996). Some of the proposed migration mechanisms for various amounts of added rocky material. should be dependent on the initial metallicity of the pro- The nature of the accreted material is relevant. To toplanetary disk. If planet migration is indeed more achieve a given level of enrichment, relatively less ma- likely to occur in a high-metallicity system, then such terial needs to be accreted if it is in the form of large planets will be easier to detect with the Doppler asteroids or terrestrial planets than if it is in the form of giant planets. If a giant planet has the same metallicity as its host star, then no enrichment is possible. However, 14While most discussion in the literature of metal enhance- the metallicity of Jupiter is about two to three times that ment of stellar atmospheres via accretion has employed the of the Sun. Gonzalez (1998a) estimated that the addition term ‘‘pollution’’ or ‘‘self-pollution,’’ we have decided to adopt of Jupiter to the Sun would increase its [Fe/H] by about the more positive term ‘‘self-enrichment.’’ 0.05 dex.

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 113

FIG. 4. Tc slope corrected for a linear trend with [Fe/H]. The data are from Smith et al., 2001; ϩ, field stars; ᭹, stars with planets.

The evidence favoring the self-enrichment explana- 6000 K. This tends to confirm the self-enrichment expla- tion is more varied. A small amount of self-enrichment nation. (ϳ0.03 dex) was proposed by Laws and Gonzalez Laughlin (2000) found that the F-type SWP’s display a (2001) to account for the Li abundance, metallicity, and higher mean metallicity than comparable field stars Tc trend differences between 16 Cyg A and B. All three without planets. This is expected from a self-enrichment differences are consistent with 16 Cyg A having accreted scenario, because the convection zone decreases rapidly more planetary material than 16 Cyg B. with increasing mass for main-sequence stars. The results of the Gratton et al. (2001) study also con- More direct evidence for self-enrichment was re- firm the self-enrichment explanation, even though none ported by Israelian et al. (2001) in the form of 6Li in the of their sample stars are known to have planets. Two of atmosphere of the SWP HD 82943. If confirmed, this the common proper-motion pairs in their study have sig- result would have provided strong evidence of self- nificantly different Fe abundances, and one pair, HD enrichment, since a star of its mass would have de- 219542 A and B, shows a differential abundance trend stroyed all its primordial 6Li by the time it reached the with Tc . This pair also display a large difference in Li main sequence. Reddy et al. (2002) did not confirm the abundances, but their relatively large temperature dif- presence of 6Li in HD 82943, but they did find a high ference, ϳ280 K, makes it likely that they have experi- upper limit for HD 219542 A. The accretion event enced significant differential Li depletion. Nevertheless, would have had to be recent in the case of HD 219542 the pattern of chemical differences in this pair is very A, given that 6Li is quickly destroyed in the atmosphere similar to that of the 16-Cyg system. of a star of its temperature. Gratton et al. (2001) derived The correlation between Tc slope and [Fe/H] reported an age of 1–2 Gyr for this pair, so it is possible that an by Smith et al. (2001) has a strong component resulting accretion event might have occurred recently. Interest- from Galactic chemical evolution. If self-enrichment is ingly, the derived relative ages for the pair HD 219542 A important, then there should also be a trend with Teff . and B, based on their relative locations on the H-R dia- This can be tested if the trend with [Fe/H] is subtracted gram, are different (see Fig. 3 of Gratton et al., 2001); from the data. The result of such an operation is shown HD 219542 A appears to be about 2 Gyr younger. This in Fig. 4. Both groups of stars in the figure display con- anomaly can be understood if the atmosphere of HD siderable scatter, but there is a noticeable clumping of 219542 A is enriched by about 0.2 dex more metals than the SWP’s at large values of the Tc slope for Teff above its companion (see discussion on 51 Peg below).

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In a very different kind of study, Murray et al. (2001) bits. One possible test could be to look for a trend with reported on metallicity trends that correlate with the Li semimajor axis only among the planets in orbit about abundances among field stars. Their method is based on the cooler SWP’s, which are not expected to display the the assumption that the pattern of Li abundances among effects of self-enrichment. dwarfs and subgiants tells us something about changes in It could very well be the case that all the extrasolar the extent of their outer convection zones and thus how giant planets known to date are ‘‘short-period’’ systems any initial Fe enrichment might have been diluted. They that have experienced substantial migration. Perhaps found that the addition of an average of 0.4 M of Fe to there is a group of systems with giant planets like Jupiter the field stars can account for the observed trends. that have not experienced substantial migration, but we Murray and Chaboyer (2002) analyzed the SWP shall have to await the discovery of such planets before sample in the same way as their field-star sample and we can determine if there is a distinct class of ‘‘long- found that the observed metallicity trends with stellar period’’ systems. mass and age are consistent with the addition of an av- erage of 6.5 M of Fe in each star. B. Evaluation Barbieri and Gratton (2002) cited Galactic kinematics of the SWP’s as evidence for self-enrichment. They base The various research groups cited above have not this on the apparent dependence of the metallicities of reached a consensus as to the source(s) of the high me- the SWP’s on their perigalactic distances (more pre- tallicities of the SWP’s. At the time of this writing, the cisely, the lower envelope of the metallicity versus peri- Laws and Gonzalez (2001) differential analysis of the galactic distance distribution is not flat). If self- 16-Cyg system, the Gratton et al. (2001) differential enrichment is not the primary source of their high analyses of other common proper-motion pairs, the metallicities, then there should be no dependence on Smith et al. (2001) Tc trends, and the Murray and Galactic location. In other words, if only the primordial Chaboyer (2002) study provide the most convincing evi- explanation is important, then giant planets will form dence for self-enrichment. [The work of Murray et al. once some critical minimum metallicity is exceeded, re- (2001) is not included in this list due to the heteroge- gardless of the Galactic kinematics. The results of neous nature of the published metallicity values they Gonzalez, Laws, et al. (2001) on Galactic positions of employed in their statistical analyses of field stars.] young SWP’s also tends to support this interpretation. Pinsonneault et al. (2001) note that the level of self- enrichment needed to account for the metallicities of solar-temperature SWP’s would lead to unrealistically high metallicities for the hotter SWP’s. Their criticisms 3. Migration of the self-enrichment explanation are helpful, but they The fact that the Doppler planet search method is do not refute it. [The reader is referred to Murray and strongly biased towards large velocity amplitudes means Chaboyer (2002) for additional discussion of their that any survey employing it will be more sensitive to work.] In addition, the extremely high metallicities pre- massive close-in planets. Therefore any physical process dicted by Pinsonneault et al. (2001) for the hotter SWP’s that leads to a planet in an end state with a large Dop- are reduced if one considers the possibility that engulfed pler amplitude will be preferentially detected. One such planets could penetrate the thin outer convection zones process is planet migration. If it is sensitive to the initial of the more massive stars (Sandquist et al., 1998, 2002). metallicity, such that more metal-rich systems are more What is more, the enrichment of a star’s atmosphere by likely to experience migration, then a Doppler survey rocky material would increase the depth of the convec- will show a correlation between the presence of planets tion zone, since its depth depends on metallicity. Thus and the metallicity of the host star. Dynamical simula- the existing calculations of stellar atmosphere enrich- tions of giant-planet migration via scattering of plan- ment need to be redone with this effect taken into ac- etesimals indicate that substantial amounts of iron can count. be dumped onto their host stars (Murray et al., 2002; Ryan (2000) provided an alternative explanation for Quillen and Holman, 2000). the different lithium abundances of the 16-Cyg pair. He One possible way to test for this bias is to search for a argued that the temperature difference between these metallicity trend with semimajor axis. Gonzalez (1998b) two stars could, over their histories, lead to sufficient and Queloz et al. (2000), working with small samples of differential depletion of lithium to account for the ob- SWP’s, observed that those with hot Jupiters tended to served difference. However, this explanation cannot ac- be more metal rich than those without. Santos et al. count for the observed difference in iron between them. (2001) revisited the question with a sample containing High-precision differential comparisons are not possible just over 40 SWP’s and did not find a clear trend. for the HD 178911 system, given that the primary is Unfortunately, because planet migration might be composed of two unresolved stars. And, even if it were closely linked with self-enrichment, it is difficult to pro- possible, it is unlikely that self-enrichment was possible pose a clear-cut test for migration. Also, it is not clear in this system given the small separation of the stars. how to treat multiple-planet systems, and some stars At the same time, some observations argue against presently known to have only one giant planet may later self-enrichment as the dominant source of the high me- be found to have additional giant planets in larger or- tallicities of the SWP’s. In particular, the high metallici-

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 Guillermo Gonzalez: Stars, planets, and metals 115 ties of the subgiants, HD 27442, HD 38529, and HD less, the results of Barbieri and Gratton (2002) and 177830, argue for the primordial explanation. In addi- Gonzalez, Laws, et al. (2001) are strongly suggestive. tion, the fact that the two most metal-rich SWP’s, 14 Her ␳1 and 55 Cnc, are late-G dwarfs is surprising within a VII. THE SOLAR SYSTEM self-enrichment framework such as that of Murray and Chaboyer, as it would predict the highest metallicities We should not forget that the Solar System is also a for mid-F SWP’s instead. One could counter by noting data point in the SWP dataset. However, there are two that late-G dwarfs are more common than mid-F dwarfs. significant differences that we cannot overlook. First, the We should then expect a very metal-rich mid-F dwarf planet-hunting groups are only now reaching the needed (near ͓Fe/H͔ϭϩ0.5) to be discovered in the near future time baseline to reliably detect the signature of a true as the sample size continues to grow. Jupiter twin (with an orbital period near 12 years). So at The detection of 6Li in HD 82943, if it had been con- this time we cannot fairly compare the properties of the firmed, would have counted strongly in favor of the self- Sun to those of the presently known SWP’s. Second, by enrichment explanation. In fact, the firm detection of 6Li virtue of being inhabitants of the Solar System, we can- in any metal-rich F dwarf (not just one with planets) not neglect observer selection bias resulting from ‘‘weak would lend strong support to the self-enrichment expla- anthropic principle’’ considerations (Barrow and Tipler, nation. 1986). In other words, some of the particular parameter Unfortunately, studies of Be abundances in SWP’s values we measure for the Sun and the orbital param- have not proved very useful in testing the self- eters of Jupiter may be required for our existence enrichment hypothesis. Be is destroyed much more (Gonzalez, 1999). Giant planets orbiting other stars are slowly than is Li in a stellar atmosphere. Thus the accre- not subject to this selection bias. tion of high-Z matter will have a relatively smaller effect Having said this, there is some insight to be gained on Be. And, since most enrichment mechanisms favor about extrasolar planetary systems by studying the Sun. early accretion rather than late, this will make it even We can apply some observational experiments to the less likely that Be could be employed to detect enrich- Sun that are not possible for other stars. One such test is ment. helioseismology (while not impossible for other stars, it We should not overlook the possibility that the prop- is far less information rich), which allows us to recon- erties of the planets themselves can help us choose struct the radial variation of sound speed in its interior. among these competing explanations. As noted above, a This places strong constraints on possible radial metal- metallicity trend with semimajor axis is indicative of licity variations resulting, for example, from a polluted both migration and self-enrichment. Murray and atmosphere. Inhomogeneous solar models were dis- Chaboyer have noted the presence of a weak correlation cussed in the literature as a possible solution to the solar between metallicity and planet minimum mass, which neutrino problem as early as 1971 [see Yang et al. (2001) they note is consistent with a class of migration models for a review of the literature]. The most recent studies resulting from planetesimal scattering. Even the planets’ indicate that such models are unlikely to solve the solar eccentricities can be informative. Israelian et al. (2001) neutrino problem by themselves (Christensen-Dalsgaard argued that the high eccentricities of the planets around and Gough, 1998; Christensen-Dalsgaard, 2002). HD 82943, when interpreted within the framework of Gonzalez (1997) presented tentative evidence for a planet-planet scattering, are consistent with planet en- trend with Tc in the Sun’s photospheric abundances rela- gulfment. A similar argument has been offered for the tive to meteorites and interpreted it as possibly resulting 16 Cyg binary system. However, even if the orbital sig- from enrichment. Figure 5 is an updated version of Fig. nature of planet-planet scattering in a given system is 5 in Gonzalez (1997). While a weak trend is visible in strong, that does not guarantee that a planet was en- the plot, the large scatter prevents us from reaching a gulfed by its host star; it could have been ejected from firm conclusion at this time. It is notable that the two the system altogether. Given the great variety of out- elements with Tc less than 1000 K in the figure are defi- comes of planetary dynamic interactions, useful con- cient in the solar photosphere. Two processes are re- straints from the planets’ properties will be possible only sponsible for the general trend in Fig. 5. One is the when we have enough systems so that we can form sta- depletion of highly volatile elements in meteorites; C, N, tistically interesting groups within the sample. It could and O (not shown in the figure) are depleted in meteor- be the case that the properties of the planets and their ites by 1.2, 1.6, and 0.3 dex relative to the solar photo- host stars will naturally clump in certain regions of pa- sphere, respectively. The other process is the enrichment rameter space. For example, the hot-Jupiter systems of the Sun’s atmosphere by high-Z material. If both pro- form a fairly distinct class. Even then, it will not be ob- cesses operated in the early Solar System, then the ex- vious how we should treat multiple-planet systems. pected trend in Fig. 5 would be for large positive devia- Galactic kinematics is an important independent test tions for elements with low Tc , reaching a minimum of the various proposals. However, due to dynamical dif- value at some intermediate Tc , and then increasing fusion of orbits combined with the Galactic metallicity again gradually for higher Tc . Thus, to estimate the gradient, it can be a messy business. For this reason, magnitude of any possible enrichment of the Sun’s atmo- large samples of both SWP’s and field stars are required sphere, it is first necessary to determine the range of Tc before a definitive conclusion can be reached. Neverthe- over which depletion in meteorites is significant. Given

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 116 Guillermo Gonzalez: Stars, planets, and metals

FIG. 5. Difference between so- lar photospheric and meteoritic abundances. Only those ele- ments with quoted uncertain- ties less than or equal to 0.06 dex in the Sun and meteorites are plotted. Most abundances are from Grevesse and Sauval (1998). The following sources were employed for solar photo- spheric abundances not taken from Grevesse and Sauval (1998): Mg, Si, Fe (Holweger, 2001); P (Berzinsh et al., 1997); Ti (Bizzari et al., 1998;) Ge (Biemont et al., 1999); La (Lawler et al., 2001); Pb (Biemont et al., 2000). The me- teoritic abundance of P is from Wolf and Palme (2001). The condensation temperatures are from Table 2.3 of Lodders and Fegley (1998).

ϳ this, elements with Tc below 1000 K in Fig. 5 might and metallicity. With these two criteria, the SWP’s HD need to be shifted to even more negative deviations. 95128, HD 186427, HD 195019A, and HD 222582 are There is a great need to improve estimates of solar most like the Sun. The presence of a stellar companion photospheric abundances for elements of low Tc , such could affect the final state of a planetary system. Re- as S, Zn, Cd, In, Sn, and Tl.15 moving the common proper-motion pairs leaves HD Jeffrey et al. (1997) examined this question from the 95128 and HD 222582 as the closest match to the Sun in vantage point of the total amount of material that might the SWP sample. HD 222582 has a 5.4-MJ planet in a have been accreted by the Sun as a result of planet for- highly eccentric orbit, while HD 95128 has two planets mation processes. Even if their estimates are too large, it in multiyear nearly circular orbits. Thus it is not enough is difficult to avoid the conclusion that some high-Z ma- that a star closely resemble the Sun for it to be accom- terial must have fallen into the Sun. Yang et al. (2001) panied by similar planets. However, it is notable that the found that an inhomogeneous solar model (with an en- planetary system around HD 95128 more closely re- velope polluted by 0.13 MJ of metals) is consistent with sembles the Solar System than any other system found everything we know about the Sun. to date. We can also approach the question of the status of the Solar System from a different angle: do the planetary VIII. IMPLICATIONS AND FUTURE RESEARCH systems around stars very similar to the Sun resemble the Solar System? The stars listed in Table I are ‘‘Sun- A. Present research like’’ in the sense that they occupy what is often called the lower main sequence. We can narrow our compari- 1. Learning about planet formation and evolution son from the parameters we list in the table. Arguably, the two most fundamental stellar parameters are mass Understanding of the correlation between the metal- licity of a star’s atmosphere and the presence of giant planets around it will eventually give us useful con- straints on planet formation and evolution processes. At 15The latest estimate for the solar photospheric abundance of this time we cannot cleanly separate the three proposed In places it 0.74 dex above the meteoritic value, suggesting that In is depleted in the latter (Bord, 2002). If this is the case, then explanations (primordial, self-enrichment, and migra- tion) for this correlation, but this task should only be- elements with Tc smaller than or equal to that of In, 470 K, will be systematically depleted in meteorites; however, the solar come easier as the sample size continues to grow. Each ϭ photospheric abundance of Cd, with Tc 430 K, agrees with its of the explanations has the potential to tell us something meteoritic abundance. important.

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If the primordial explanation is found to be signifi- Doppler searches continues to increase, the brown cant, then this could allow us to select between the core dwarf desert will begin to disappear. instability accretion model (Podolak et al., 1993) and the gravitational instability (Boss, 1997) models (Gonzalez, 1998a). A high degree of sensitivity to the primordial metallicity is expected to favor the core instability accre- 3. Implications for astrophysics tion model. Boss (2002) argues that the gravitational in- stability model should have very little dependence on Detection of a substantial self-enrichment signal metallicity. Perhaps both mechanisms operate, with core would have implications for the solar neutrino problem, instability accretion becoming dominant for stars with galactic chemical evolution models, and the derivation solar metallicity or above. of stellar ages (Gonzalez, 1997). One of the most basic Any detectable contribution from self-enrichment has underlying assumptions in stellar astrophysics is that the the potential to tell us a great deal about planet forma- surface metallicity of a late-type dwarf star (F-M spec- tion processes. Because the convection zone of a Sun- tral type) is an accurate indication of its overall metal- like star shrinks rapidly upon reaching the main se- licity. If the outer convection zones of F and early G ϳ quence, the amount of the observed enrichment will be dwarf SWP’s are enriched by 0.2 dex or more, then the strongly dependent on its timing [see Fig. 1 of Murray metallicity distribution constructed from observations of et al. (2001) and Fig. 1 of Laughlin and Adams (1997)]. nearby stars will be significantly in error. This problem The relative enhancements of the elements will tell us could be greatly mitigated if only dwarfs with less than ϳ about the composition of the accreted material, which, 0.05 dex worth of enrichment are employed in such in turn, may allow us to determine its identity (terres- studies. Until the quantitative contribution from self- trial versus gas giant). If gas-giant accretion is implicated enrichment is determined, it would be a good idea to in some cases, then this information may help us to limit studies of the local metallicity distribution to mid- settle the current debate about the overall metallicity of to late-G-type dwarfs. Jupiter. It is almost axiomatic that some accretion of Ford et al. (1999) calculated that if the interior of 51 high-Z material onto a star will occur if it is accompa- Peg is 0.2 dex more metal poor than its surface, then nied by a disk containing condensed particles. Thus even the derived age will be increased by 3 Gyr and the an upper limit on the amount of self-enrichment would mass will be reduced by 0.11 M᭪ relative to a homo- place useful constraints on planet formation and evolu- geneous model. While this is a significant effect for tion models. stars comparable in metallicity to the Sun, it is un- Trends of metallicity with the semimajor axes of the likely to alter globular cluster ages. Globular clusters giant planets will place important contraints on migra- are probably too metal poor and too crowded to tion models. Detection of a strong migration signal form planets. However, self-enrichment will have a among the hotter SWP’s would tend to support planet significant effect on reconstructions of the age- migration models that involve planetesimal scattering metallicity relation for disk stars, as both the observed (e.g., Quillen and Holman, 2000). However, the cooler metallicity and the derived age of an enriched star will SWP’s will provide the least ambiguous signal of planet be in error. migration.

2. Distinguishing planets from brown dwarfs 4. Implications for astrobiology The rarity of substellar-mass objects with minimum Some of the possible explanations for the correlation masses greater than about 13 MJ (brown dwarfs) as close between metallicity and the presence of planets dis- companions to Sun-like stars has made it difficult to cussed above should have implications for astrobiology study them in a statistically meaningful way. Neverthe- research. For example, the possible dependence of less, a few brown dwarfs are being found with the Dop- giant-planet migration on metallicity would have signifi- pler search method. Eventually, there will be enough de- cant consequences on the formation of terrestrial plan- tections to begin to answer some questions. In particular, ets close to their host stars. The possible dependence of it would be interesting to compare the metallicity distri- giant-planet mass on metallicity would also have an ef- butions of brown dwarf and giant-planet host stars. One fect on the stability of smaller planets in a given system. might expect a difference if, for example, brown dwarfs The mere presence of giant planets, too, can be benefi- and giant planets form via different mechanisms. Do all cial to habitability (see, for example, Wetherill, 1994). brown dwarfs form like stars, or do some form like giant Terrestrial planet formation should be even more planets? Do brown dwarfs migrate more than giant strongly dependent on metallicity than that of gas giants, planets? Answers to these questions might help us un- given the fact that they consist almost entirely of metals. derstand why brown dwarfs are common in the field Gonzalez, Brownlee, and Ward (2001) and Lineweaver (Chabrier, 2002) but very rare in close proximity to Sun- (2001) consider the connections between metallicity and like stars. Since this distribution implies that they will be habitable planets within the context of chemical evolu- found with increasing frequency beyond some minimum tion in the Milky Way Galaxy and the broader universe. orbital period, it is likely that as the time baseline of the Empirical data on the frequency of terrestrial planets

Rev. Mod. Phys., Vol. 75, No. 1, January 2003 118 Guillermo Gonzalez: Stars, planets, and metals may soon come by way of the Kepler mission, which will parsecs would be ideal. All the stars in such a sample search for the photometric signatures from transiting should receive the same quality of analysis as the SWP’s. planets.16 Only when a high quality field control sample is success- fully procured will we be able to search for many of the B. Recommendations for future observational research trends discussed in this colloquium and correct for the various biases. There is still much to be learned about the relation- ships between stars and planets. Only one trend is firmly ACKNOWLEDGMENTS established at the present time—the high mean metallic- ity of SWP’s relative to field stars. However, it is still not The author thanks Chris Laws for helpful comments clear how much each of the three proposed causes con- and Verne Smith for his data on condensation tempera- tribute to the difference. Spectroscopic observations of ture slopes and permission to include his figure herein. stars in open clusters and common proper-motion pairs The comments and suggestions of Anthony Starace led arguably have the greatest potential towards resolving to substantial improvements in the clarity of this review. this issue. The 16 Cyg and HD 219542 pairs are sugges- tive, but more systems like them need to be studied. REFERENCES Particularly helpful would be the discovery of planets around stars in the metal-rich Hyades cluster (at present Barbieri, M., and R. G. Gratton, 2002, Astron. Astrophys. 384, being surveyed by several groups). Such studies should 879. be followed up with differential spectroscopic abun- Barrow, J. D., and F. J. 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