Statistical Properties of Brown Dwarf Companions: Implications for Different Formation Mechanisms 3

arXiv:1303.6442v1 [astro-ph.EP] 26 Mar 2013 ⋆ 0.7% of rate Californi occurrence The an technique. measured RV Search by Planet far hun ﬁve Carnegie Sahlmann so over 2008; detected to compared Metchev planets therein) & giant references (Reid and known 2011a al. are et dozens few a only

o.Nt .Ato.Soc. Astron. R. Not. Mon. 0Spebr2018 September 10 ls Dcmaino oa yesa swl ihntedete the ( within surveys well RV precision is high star the type a of solar by sensitivities a tion signal RV on reflex companion induced BD the close Although 2006). Greth Lineweaver 2000; Butler & & (Marcy stars withi FGKM planets main-sequence des to around dwarf relative brown companions a dwarf of brown identification of paucity the is surve decades majo (RV) two the past velocity of the One radial 1995). high-precision Queloz & of came (Mayor achievements b 1996) Peg al. planet 51 extra-solar star, et first type the Rebolo solar of Oppen- 1996; discovery 1995; the al. as al. year et same Basri et the 1995; Nakajima al. al. cores 1995; et et Burrows al. inner o heimer et discovery their first (Rebolo 2000; The BD in 2011). Baraffe bona-fide Milsom & hydrogen & Chabrier Burrows, burn Spiegel, 2001; 1997; to al. enough et (Burrows not but terium rw wrs(D r ntems ag fapproximately of range mass the in are 13 (BD) dwarfs Brown INTRODUCTION 1 oMa Bo Mechanisms Implica Formation Companions: Different Dwarf Brown of Properties Statistical 1 c eateto srnm,Uiest fFoia 1 Bryant 211 Florida, of University Astronomy, of Department E-mail:[email protected]fl.edu − 00RAS 0000 0Jptrmse,hvn ufiin asst undeu- burn to masses sufficient having masses, Jupiter 80 1 ⋆ n inGe Jian and 000 0–0 00)Pitd1 etme 08(NL (MN 2018 September 10 Printed (0000) 000–000 , 1 e words: Key r possible their formation. and mechanisms formation dwarf brown discove about Our fragmentation. pro cloud the molecular g through fu in this binary above form Our dwarfs primarily brown origins. may while different instability, gap gravitational this have below may dwarfs they brown that that implies m a also medium which dwarfs but and Brown distribution, surveys. period eccentricity RV short different previous its significantly by detected to easily due be gap des can dwarf this brown cause the to in unlikely land” “driest the representing days), ABSTRACT ftebondafdsr.Tedt lal hwasotperio short ( a diagram show distribution clearly period-mass data dwarf The lite brown desert. from the stars dwarf host brown of their the all and of collect stars we type orbitin paper solar dwarfs around this brown dwarfs In of desert. paucity dwarf brown the so-called to due understood poorly h asdmi hr asv xrslrpaesadbond brown and planets extrasolar massive where domain mass The ∼ tr:bondaf–tcnqe ailvelocity radial technique: – dwarf brown stars: 3 − pc cec etr ansil,F,36125,USA 32611-2055, FL, Gainesville, Center, Science Space ruda around 10 sover ys AU 3 n m/s), r,a ert, dred & a a f ± c- er in r r opnosmr asv hnJptr 11 Jupiter: than massive more companions yr) hc otis10 oa yesaswti 0pc. 50 within stars type sea solar planet CORALIE 1600 the contains of which sample stellar uniform the on based ino h opnosaon erySnlk tr.Te fo investiga They stars. detailed Sun-like nearby approximately around a that companions the performed of tion (2006) Lineweaver & Grether tr Wteme ta.20) alane l 21a obta (2011a) al. et of limit Sahlmann upper 2009). an al. et (Wittenmyer stars ae ta.20) n h coadOsraoyPae Sear Planet 0.8% Observatory of McDonald rate similar the a and shows 2007), al. et Patel .%fo hi apeof sample their from 0.2% 1 fBsi 28 3 in freq the BDs inferred (2009) of Hillenbrand & Metchev stars. young 25 seas ece ilnrn 04 sa ls separation close a at obtained as separati (2007) 2004) wide at Hillenbrand al. Lafrenière rare & et Metchev as be also not might (see Gizi BDs stars, type that solar suggests around (2001) rare are BDs close the Although companions, rud26Snlk stars. Sun-like 266 around . . 9 2 − +8 − +3 − 5A nteGmn epPae uvyaon 5nearby 85 around Survey Planet Deep Gemini the in 250AU 1 2 oass h elt ftebondafdesert, dwarf brown the of reality the assess To D r rdtoal eivdt omlk tr,through stars, like form to believed traditionally are BDs . . . . 3 1 5 7 % % A T E rma dpieotc uvyfrsbtla companions substellar for survey optics adaptive an from o h rqec f13 of frequency the for tl l v2.2) file style X < − 50A risaon on oa nlg is analogs solar young around orbits AU 1590 1% 0 . 16% 6% 35 r D,and BDs, are o nycnrsta hsgpi real, is gap this that confirms only not o h rqec fcoeB companions BD close of frequency the for fnab u-iesashv ls ( close have stars Sun-like nearby of m < pmydmnnl omlk stellar a like form dominantly may ap ishv fee motn insights important offered have ries ∼ aueadsuytedemographics the study and rature ± r.Osrainbae r highly are biases Observation ert. vial aaaotcoebrown close about data available 1000 ltosiswt lntadstar and planet with elationships sin .%fo erhsml f250 of sample search a from 0.6% te ttsia td indicates study statistical rther ls oslrtp tr,the stars, solar-type to close g oeadblwti a have gap this below and bove n eimms a in gap mass medium a and d olntr iktruhdisk through disk toplanetary − < i s,o hc rw dwarfs brown which of ass, 75 5% agtsas(ote l 2002; al. et (Vogt stars target 95% M 55 Jup ± af vra sstill is overlap warfs ofiec nevlof interval confidence M 2% opnosbetween companions Jup % r in planets. giant are in for tions ± and 3% < P r stellar are < P tal. et s uency ined 100 rch, und ons ch s. 5 - 2 Bo Ma1E-mail:[email protected]fl.edu and Jian Ge11Department of Astronomy, University of Florida, 211 Bryant Space Science Center, Gainesville, FL, 32611-2055, USA gravitational collapse and/or fragmentation of molecular clouds (Padoan & Nordlund 2004; Hennebelle & Chabrier 2008). A recently found self-gravitating clump of gas and dust has a mass (0.015 − 0.03 M⊙) in the BD regime (André, Ward-Thompson, & Greaves 2012), which supports the idea that BDs could form like a star. On the other hand, companions with masses up to 10 MJup (Alibert et al. 2005) or even 38 MJup (Mordasini et al. 2009) may form in protoplanetary disks according to the standard core-accretion planet formation theory. Because of this, the brown dwarf desert is commonly interpreted as the gap between the largest mass objects that can be formed in protoplanetary disks, and the smallest mass clumps that can collapse and/or fragment in the vicinity of a protostar. The mass function of close stellar companions shows a linear decrease in log(M) toward the BD mass range from both stellar mass and planetary mass directions (Grether & Lineweaver 2006). In comparison, the mass function of isolated substellar objects Figure 1. Period distribution of known BD companions around solar type 20 M seems to be roughly flat in log(M) down to masses ∼ Jup, stars. both in the field and in clusters (Luhman et al. 2000; Chabrier 2002). This indicates that close BD companions may form in a different way from those formed in the field and clusters. As such, statistical properties of close BD companions as well as how these properties are related to their host stars, contain a lot of information about the poorly understood BD formation mechanisms and their relationships with star and planet formations in close orbital environments. These statistics may also be important to investigating how additional important processes such as tidal evolution and disk-planet interaction affect close BD properties (e.g., Armitage & Bonnell 2002; Matzner & Levin 2005). Given that the close BD occurrence rate is < 1%, currently there has yet been a single large, relatively uniform RV survey capable of produc- ing a large homogeneous sample of BD companions for a meaning- ful statistical study (Marcy et al. 2000; Ge et al. 2008; Sahlmann et al. 2011a). However, all of the previous RV planet surveys have sufficient RV sensitivity and time baseline (at least more than 2 years, e.g., Ge et al. 2008; Ge et al. 2009; Eisenstein et al. 2011) to detect a majority of close BDs around solar type stars due to the large Figure 2. Cumulative mass distribution of brown dwarf candidates. Three RV amplitude (on the order of ∼ 1000 m/s, vs. a few to a few tens lines with three RV precisions, 50 m/s, 100 m/s and 150 m/s, are also shown. m/s RV precisions in previous RV planet surveys). It is, therefore, possible to combine close BD companion samples together for a statistical study without major biases. For those which are transiting their parent stars or have astrometry In this paper we assembled a catalog of all the BD companions measurements, true masses are also given besides M sin i. discovered around solar type star from literature and used it to con- duct a statistical study. We have found tentative evidence for the ex- istence of two different populations of BD companions. We present 3 OBSERVED PROPERTIES OF BROWN DWARFS the BD catalog assembled in this study in §2 , statistical properties of BD companions in §3 and discuss their implication for different 3.1 Orbital Period Distribution BD formation scenarios in §4. We summarize our main results in The distribution of orbital periods of BDs has two main features §5. (Fig. 1): a relatively flat distribution inside P ∼ 100 days and a sharp jump beyond P ∼ 100 days. The drop beyond P ∼ 1000 days is likely due to the observational incompleteness since 1) it is more difficult to detect a BD companion over a long period than a short period with RVs and 2) some RV surveys (such 2 CATALOG DESCRIPTION as the SDSS-III MARVELS, Ge et al. 2008; Ge et al. 2009; We have collected data from literature about the currently known Eisenstein et al. 2011) do not cover beyond this period. It is ev- BD (candidates) companions around FGK type stars. Most of them ident that the number of BDs increases with the orbital period, have known Keplerian orbits, except for HIP 78530 (discovered even though RV and transit observations are biased toward dis- by direct imaging; Lafrenière et al. 2011), KOI-205.01 (unpub- covering objects in short-periods. The position of the maximum lished yet; Santerne et al. 2012) and KOI-554.01 (unpublished yet; of the distribution is unknown due to the different duration limit Santerne et al. 2012). Properties of the BDs (minimum mass, pe- of most of the old surveys (several thousand days). This increas- riod and eccentricity) and their host stars (mass, effective temper- ing distribution is consistent with the results from high contrast ature, surface gravity and metallicity) are summarized in Table 1. and high angular resolution imaging surveys, which find evidence

c 0000 RAS, MNRAS 000, 000–000 Statistical Properties of Brown Dwarf Companions: Implications for Different Formation Mechanisms 3

lative histogram for the long period group. We carried out a sim- ple Monte Carlo experiment to test the emptiness of this gap on the period-mass diagram. There are a total of 25 BDs with period shorter than 100 days in our BD sample. We assumed simply that their masses are uniformly distributed between 13 and 80MJup. Then we drew their masses randomly from this uniform distribution and counted how many of them will fall in the mass range of 35 to 55MJup. We found the probability that less than 3 BDs will fall in this gap is 0.9%, which corresponds to a 2.6σ significance. A larger BD sample in the future will be better to assess the significance of this gap. The appearance of this depleted region in the BD period-mass diagram has naturally divided BDs into two mass groups: one with masses greater than 42.5MJup and the other less than 42.5MJup. Their properties and origins may be different. We further explore properties of these two groups and study possible origins. Figure 3. Cumulative mass distribution of brown dwarf companions. Brown dwarfs with periods greater and less than 100 days are shown as dashed and 3.3 Orbital Eccentricity Distribution solid lines respectively. The orbital eccentricities show great difference for the two BD groups with masses greater and lower than 42.5MJup, respectively. for a higher fraction of BD companions at wide orbits than at Fig. 4 shows the period-eccentricity distribution of all known BDs. close orbits (Lafrenière et al. 2007; Metchev & Hillenbrand 2009; The period-eccentricity distribution of BDs with masses greater Chauvin et al. 2010; Janson et al. 2012). than 42.5MJup is consistent with a circularization limit of ∼ 12 For comparison, both extrasolar giant planets days, which is similar to that found in nearby stellar binaries (Cumming, Marcy, & Butler 1999; Udry, Mayor, & Santos (Raghavan et al. 2010). It is clear that there are a significant num- 2003; Marcy et al. 2005; Udry & Santos 2007) and binaries ber of BDs with 300 d

day pile up with the than 42.5MJup. We have done a two-dimensional Kolmogorov- BD distribution as shown in giant planets (Udry et al. 2003). The Smirnov (K-S) test for the period-eccentricity distribution of BDs reason may be that BDs initially form in the protoplanetary disks with masses greater and lower than 42.5MJup. The probability that further away from the host stars where protoplanetary disks have these two BD samples are drawn from the same distribution in the more materials to efﬁciently form BDs than at the close-in regions, period-eccentricity plane is 1.7%. and the migration mechanism may not efﬁciently move such a Next we are going to compare the period-eccentricity distribu- massive body to a short period orbit (Trilling et al. 1998; Nelson tion of the BDs to that of stellar binaries. Halbwachs et al. (2003) et al. 2000; Trilling, Lunine & Benz 2002). On the other hand, have studied the statistical properties of a sample of 89 FGK type BDs forming at the same time as the primary stars may migrate main-sequence binaries with periods up to 10 year. Here we choose inward quickly in the initial gas rich disks and thus be destroyed to use their 89 binary sample for our comparison. To compare the via mergers with the stars (Armitage & Bonnell 2002). period-eccentricity distribution between BDs and stellar binaries we made use of a two-dimensional K-S test. The probability for the period-eccentricity distribution to be the same is 18% between BDs 3.2 The Period-Mass Diagram with masses above 42.5MJup and the stellar binary sample, and The period-mass distribution of BDs shows a statistically signif- 0.1% between BDs with masses below 42.5MJup and the stellar icant gap at the short period and medium mass region as illus- binary sample. These results suggest that BDs with masses greater trated in Fig. 2. In this plot, a rectangular area with P < 100 d than 42.5MJup have a very similar period-eccentricity distribution and 35MJup 100 days) and plotted their mass cumulative histograms in Fig. 3 for 3.4 Metallicity of the BD Host Stars comparison. It is apparent that a depletion of BDs with masses between 35 and 55MJup appears in the cumulative histogram for the The BD host stars in this study have a mean metallicity of [Fe/H] short period group while no depletion of BDs appears in the cumu- = −0.04 with a standard deviation of 0.28. For comparison,

c 0000 RAS, MNRAS 000, 000–000 4 Bo Ma1E-mail:[email protected]ﬂ.edu and Jian Ge11Department of Astronomy, University of Florida, 211 Bryant Space Science Center, Gainesville, FL, 32611-2055, USA

Figure 4. Period-eccentricity distribution of brown dwarf candidates. Brown Figure 5. Mass-eccentricity distribution of exoplanets and brown dwarf dwarf candidates with masses above and below 42.5MJup are shown as candidates. The solid line shows a trend that the more massive the gi- circles and diamonds. Brown dwarf candidates with true masses measured ant planet/BD is, the smaller maximum eccentricity it tends to have, using transiting observations or astrometry measurements are shown as which breaks at ∼ 42.5MJup. The vertical dotted line shows M sin i = crosses. 42.5MJup. Planets are shown as squares. Brown dwarfs with masses above and below 42.5MJup are shown as triangles and circles respectively.

Raghavan et al. (2010) has carefully studied a sample of 454 nearby solar type stars, which have a mean metallicity of −0.14 with a standard deviation of 0.25. Six stars in that sample have BD companions, which have a mean metallicity of [Fe/H]= −0.05. The mean metallicity of our BD sample is slightly higher than the mean metallically of volume limited nearby FGK dwarf stars. For ex- ample, Favata, Micela, & Sciortino (1997) have analyzed a volume limited sample of 91 G and K dwarfs, yielding a mean metallicity of [Fe/H]= −0.08 with a standard deviation of 0.26. Nordström et al. (2004) have derived metallicity for 16682 nearby F and G dwarf stars with a mean of −0.14 and a dispersion of 0.19 dex. Sousa et al. (2011) found that a mean metallicity for the CORALIE survey sample of 1248 stars and the HARPS survey sample of 582 stars is [Fe/H]=−0.11 and −0.10, respectively. However, since some exoplanet surveys choose samples biased towards metal rich stars (e.g., Valenti & Fischer 2005), the slightly higher mean metallicityof BD host stars is possibly caused by the sample bias. After comparing the BD host star metallicities with the vol- Figure 6. Cumulative metallicity distribution of host stars with BD com- 42.5M ume limited sample from Sousa et al. (2011) and the planet search panion masses above (dashed line) and below (solid line) Jup. Also sample from Valenti & Fischer (2005), we find that the BD host shown for comparison are cumulative metallicity distribution of the planet star metallicity distribution is consistent with the combination of search sample from Valenti & Fischer (2005), the CORALIE planet search sample (Sousa et al. 2011) and giant planet sample from the Exoplanet Orbit these two samples. This means we cannot interpret the BD host star Database (Wright et al. 2011). sample as metal rich. We compared metallicities of BD host stars with that of giant planet (1MJup < m sin i < 5MJup) host stars. The data of giant planet host stars is taken from the Exoplanet Or- greater than 42.5MJup have been found around stars with [Fe/H]< bit Database (Wright et al. 2011). A K-S test shows that the prob- −0.5, while several BDs with masses lower than 42.5MJup have ability of these two samples selected from the same distribution is been found in this metallicity regime. A K-S test has been con- 2 10−4 × . This indicates that the two samples are significantly dif- ducted to show that the probability of these two sample metallicities ferent from each other. are drawn from the same distribution is 70%, which suggests there We investigated the correlation between BD host star metallic- is no significant difference between these two samples regarding ities and BD masses. The Spearman’s rank correlation coefficient their metallicity distribution. of the BD mass and their host star metallicity is 0.07 with a 61% significance, suggesting there is no significant correlation between the BD mass and their host star metallicity. We also divided the BD companions into two subsample ac- 4 DISCUSSION cording to their masses. The cumulative metallicity distribution for 4.1 Two Different Brown Dwarf Populations host stars with BD companion masses above and below 42.5MJup is shown in Fig. 6. The main difference between these two distribu- Our study suggests that BDs with masses lower than ∼ 43MJup tion is at the lower metallicity end. Currently no BDs with masses have an eccentricity distribution consistent with that of giant plan-

c 0000 RAS, MNRAS 000, 000–000 Statistical Properties of Brown Dwarf Companions: Implications for Different Formation Mechanisms 5 ets in the mass-eccentricity diagram while BDs with masses above rate (Raghavan et al. 2010), i.e., lower-metallicity clouds might be ∼ 43MJup have the star-like eccentricity distribution (Figure 5). more likely to fragment to form binary stars. If this applies to mas- Our mass limit is consistent with the minimum of the mass func- sive BDs, then we expect this anti-correlation. However, our current tions of planet and stellar companions in the BD mass region, sample is too small to tell, but is at least not inconsistent with this 43+14 M −23 Jup, derived by Grether & Lineweaver (2006) using their statistics. stellar companion and giant planet sample within 50 pc around the sun. This mass function minimum is also consistent with that derived by Sahlmann et al. (2011b), who found a void in the 4.3 Upper Mass limit of Planets M mass range between 25 and 45 Jup using the data from the Our study appears to help to address the upper mass limit of ”plan- CORALIE radial-velocity survey. They further suggested that there ets”. Here we define a ”planet” as an object formed in a protoplan- may be a possible dividing line between massive planets and sub- etary disk from the companion formation point of view. It has been stellar companions. Schneider et al. (2011) has chosen arbitrarily long known that the low mass end of brown dwarfs overlaps the M and probably provisionally 25 Jup as the upper limit of mas- high mass end of massive planets formed in protoplanetary disks. sive planets based on previous studies (e.g., Sahlmann et al. 2011b; From theoretical simulations, planets with masses up to 10 MJup Baraffe, Chabrier, & Barman 2010). (Alibert et al. 2005) or even 38 MJup (Mordasini et al. 2009) may Our BD sample is, therefore, naturally divided into two dif- form in protoplanetary disks. However, it is difficult to determine 42.5M ferent groups with the mass limit of ∼ Jup. The eccentric- the upper mass limit of planet from observations since the high ity distribution of low mass BDs appears to be consistent with the mass planets detected may actually be brown dwarfs formed as low- prediction from the “planet-planet scattering” model (Rasio & Ford mass binaries. One way to tackle this problem is to search for de- 1996; Ford & Rasio 2008; Chatterjee et al. 2008), while the eccen- bris disk around massive planet/BD host star, which support the tricity distribution of massive BDs appears to be similar to that of idea that these objects are formed like a planet. Moro-Mart´ın et al. stellar binaries (Halbwachs et al. 2003) as shown in § 3.3. The ec- (2010) have found evidence of debris disk around two massive centricity distributions of these two groups support that BDs may planet host star (HD 38529 and HD 202206), which supports for- form differently: BDs below this mass limit form in protoplanetary mation of planets up to 17 MJup in a protoplanetary disk. The other disks around host stars, and above this mass limit form like stellar way to resolve this problem is through the study of mass spectrum binary systems. This is supported by our analysis results. The ex- of massive planets/BD. Motivated by the observed mass distribu- istence of a large population of long period low eccentricity BDs +14 tion, this planet upper mass limit has been set at around 43− MJup (P > 300 days and e < 0.4) serves as evidence to support the BD 23 (Grether & Lineweaver 2006) or 25 MJup (Sahlmann et al. 2011b; formation scenario in the protoplanetary disks for those compan- Schneider et al. 2011). Analysis conducted in this paper was also M ions with masses below 42.5 Jup. While the lack of long period able to provide a clue about this upper mass limit. The cumula- M and low eccentric BD companions with masses above 42.5 Jup tive BD mass distribution (Fig. 3) suggests this limit is in the range appears to support the BD formation scenario like stellar binary 30 − 60 MJup and mass-eccentricity relation (Fig. 5) suggests formation. Nevertheless, a small number of BDs in each of these that this limit is around 43 MJup. It appears that the upper limit two mass regions may form in an opposite formation mechanism, is around 43 MJup, which is consistent with the result found by but our sample is not sufficient to distinguish these minor groups. Grether & Lineweaver (2006), but is slightly larger than that from (Sahlmann et al. 2011b). 4.2 BD formation mechanisms Our study of metallicity distribution of BDs appears to support 5 SUMMARY the disk instability mechanism (Boss 1997) for those BDs formed in the planetary disks, while it is inconsistent with the predic- We have searched literature and presented a catalog of BD compan- tion of core-accretion formation mechanism (Pollack et al. 1996; ions around solar type stars found by radial velocity, transiting and Ida & Lin 2004; Alibert et al. 2005). Currently, there are two hotly astrometry observations. We have studied distribution of different contested theories about giant planet formations: core-accretion parameters of BD companions around solar type star, and found and disk gravitational instability. In the core-accretion scenario, that: giant planets form more efficiently around metal rich planetary disks than metal poor planetary disks because the higher the grain (1) BD companions have an increasing distribution with pe- content of the disk is, the easier the metal core is to build for riod, similar to giant planets and low-mass binaries; giant planets (Pollack et al. 1996; Ida & Lin 2004; Alibert et al. (2) BD companions are almost depleted at P < 100 d and 2005). In contrast, the disk gravitational instability process al- 30MJup Galaxy, followed by a roughly flat Ida & Lin 2004; Alibert et al. 2005). However, our result, showing distribution; no correlation between the BD occurrence rate and metallicity, ap- (5) Host stars of BD companions are not metal rich, and have pears to support the disk instability mechanism for formation of significantly different metallicity distribution when comparing with BDs in planetary disks. host stars of giant planets, suggesting a formation scenario at least For stellar companions, there appears to be a weak anti- partly different from the core-accretion scenario. correlation between metallicity and stellar companion occurrence The distribution of BD and their host star properties presented

c 0000 RAS, MNRAS 000, 000–000 6 Bo Ma1E-mail:[email protected]fl.edu and Jian Ge11Department of Astronomy, University of Florida, 211 Bryant Space Science Center, Gainesville, FL, 32611-2055, USA in this paper may lend support to such a picture: (1) BD compan- Fischer, D. A., Marcy, G. W., Butler, R. P., et al. 2002, PASP, 114, ions with masses below 42.5MJup form in a protoplanetary disk 529 through the the instability-fragmentation scenario, and their eccen- Fleming, S. W., Ge, J., Mahadevan, S., et al. 2010, ApJ, 718, 1186 tricity is excited through scattering with other objects formed in this Fleming, S. W., Ge, J., Barnes, R., et al. 2012, arXiv:1206.5514 disk or interactions with disk/third body; (2) BD companions with Ford E. B., Rasio F. A., 2008, ApJ, 686, 621 masses above 42.5MJup form like stars, through molecular cloud Frink, S., Mitchell, D. S., Quirrenbach, A., et al. 2002, ApJ, 576, fragmentation, similar to the formation of a stellar binary system. 478 Galland, F., Lagrange, A.-M., Udry, S., et al. 2006, A&A., 452, 709 Ge J., et al., 2008, ASPC, 398, 449 ACKNOWLEDGMENTS Ge J., et al., 2009, SPIE, 7440, 18 Gerbaldi, M., Faraggiana, R., & Caffau, E. 2007, A&A., 472, 241 This research has made use of the Exoplanet Orbit Database and the Gizis J. E., Kirkpatrick J. D., Burgasser A., Reid I. N., Monet Exoplanet Data Explorer at exoplanets.org. We thank Neil Thomas D. G., Liebert J., Wilson J. C., 2001, ApJ, 551, L163 for proof reading our paper. We acknowledg the support from DoD Grether D., Lineweaver C. H., 2006, ApJ, 640, 1051 Coopeative Agreement W911NF-09-2-0017, Dharma Endowment Halbwachs, J. L., Arenou, F., Mayor, M., Udry, S., & Queloz, D. Foundation and the University of Florida. 2000, A&A., 355, 581 Halbwachs J. L., Mayor M., Udry S., Arenou F., 2003, A&A, 397, 159 Hatzes, A. P., Guenther, E. 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c 0000 RAS, MNRAS 000, 000–000 8 Bo Ma1E-mail:[email protected]ﬂ.edu and Jian Ge11Department of Astronomy, University of Florida, 211 Bryant Space Science Center, Gainesville, FL, 32611-2055, USA 18 18 63 1 1 1 10, 11, 18, 1 32 4 1 1 58 1, 56 62 4 11, 34, 43 1 18 48, 53 36, 37 29, 52 3 11, 33, 35 1, 12, 28 4 4 32 16 1, 24, 25 45 40, 41, 42 17 46 8 18 4 21 7 38 44 31 1, 26 30 4 36, 37 12 9, 56 1, 27 1, 12, 22 5 6 05 06 12 06 05 06 07 10 06 06 05 03 06 06 02 18 10 07 04 10 07 09 06 06 10 11 02 03 05 08 . 08 ...... 05 05 13 06 05 10 07 10 07 10 01 03 18 07 . 0 09 10 10 ...... 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 . . . 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 16 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 0 ± ± ± [Fe/H] References 10 37 06 21 01 00 04 05 08 04 65 33 04 27 13 15 00 17 15 63 01 23 35 25 24 32 10 27 31 − 2 0 0 067 ...... 04 13 38 22 01 21 10 03 37 07 32 19 04 29 . . . . 430 ...... 0 0 0 0 0 1 0 0 0 0 0 0 0 0 0 0 1 0 0 0 0 0 0 . 0 − − − − − − − − − − − − − − − − − − − − − − − 1 0 1 0 03 02 15 0 10 10 11 0 10 04 0 05 12 0 10 10 10 14 04 +0 06 15 0 10 10 0 03 10 0 06 06 +0 35 20 0 10 25 +0 10 0 05 20 0 10 0 07 10 0 04 0 06 0 10 10 +0 17 0 16 10 10 10 12 +0 10 +0 20 +0 04 05 ) ...... g 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 ± ± +0 − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ··· ··· log( 4.4 2.0 4.13 4.48 4.60 4.55 4.56 4.48 4.47 4.49 4.02 1.51 2.98 4.55 4.44 4.49 4.52 4.39 4.54 4.56 4.52 4.43 4.45 4.37 4.28 4.47 2.08 4.37 4.08 4.76 2.74 1.51 4.75 4.57 3.89 4.83 4.13 4.12 4.32 4.44 2.40 4.07 2.31 4.50 4.02 4.10 4.48 29 33 44 69 10 92 40 44 42 80 24 50 50 50 65 50 50 10 50 50 44 92 50 50 26 32 44 89 80 40 90 70 21 40 40 87 82 500 110 100 100 100 100 100 100 100 103 100 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± eff ± ± ± ± ± ± ± ± ± ··· ··· ··· ··· ··· ··· 7250 4.34 5697 3.94 T -type stars 5664 5027 5223 5316 5216 5188 6131 4137 6224 5715 5833 5334 6047 5104 5491 6297 5629 4962 4830 5640 5554 5886 5555 5185 5821 6004 6427 5950 4569 4137 6186 6135 5680 5903 5700 4917 5183 6074 5000 5852 6048 5083 4742 4500 10500 08 08 09 3 03 04 02 02 01 02 23 5 04 11 2 5 3 03 04 04 02 04 1 03 10 05 09 04 1 2 021 08 040 0 03 05 4 3 05 087 ...... 018 ...... 10 10 02 05 02 026 ) (K) (cgs) . . . . . 0 0 0 0 0 0 0 0 0 0 0 1 0 0 0 1 0 0 0 0 0 0 0 0 0 0 0 0 . 0 0 0 0 0 0 0 0 0 0 . 0 . 0 0 0 ⋆ 0 1 ⊙ 78 19 75 98 . +0 − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . . +0 − +0 − +0 − 1 1.15 5765 0.28 3500 0.83 5360 M M 0 0 9 8 5 8 6 0 7 4 0 9 7 ( ...... 11 21 89 75 92 81 81 83 83 80 36 37 14 94 05 89 96 96 21 78 25 40 16 10 48 01 914 3 2 1 1 ...... 1.22 1.28 780 998 . 1.112 0 0 0 0 0 0 1 . . 0 06 2 03 2 005 1 081 1 002 0 01 1 002 0 001 095 1 005 2 01 02 1 067 . . . . 068 ...... 02 1 10 0 06 7 07 2 04 013 0 007 001 0 011 1 010 2 003 0 012 0 011 1 002 1 007 3 035 0040 1 0010 1 039 0023 0 0017 0003 0 0039 1 0073 1 0054 0010 0 0020 1 . 0019 . . 001 ...... 0 0 0 0 0 0 0 0 0 0 0 ...... 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 0 +0 − ± ± ± ± ± ± ± ± ± ± ± ± +0 − ± ± ± ± ± +0 − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 0.32 0.25 0.19 0.54 0.20 ± 0.323 0.741 0.688 0.703 0.385 0.3325 0 (ﬁxed) 18 15 23 63 685 . . . . 0.40 0.14 0.27 0.21 0.50 340 838 448 277 433 356 231 . 0.638 0.26 0.900 ...... 0.723 0.099 0.359 0.015 0.257 0.084 0.081 0.226 0.362 0.082 0 0.0017 0.4375 0.0066 0.8191 0.6780 0.7124 0.3985 0 0.6616 0.4751 0.2122 2 6 1 6 74 6 . . . . . 2 0 3 6 5 20 0 9 0 29 9 . 0 . 03 0 . 1 0 2 0 6 0 3 6 6 0 14 01 0 45 62 22 35 32 0 42 0 000 089 000 ...... 10 . 01 01 05 0 ...... 0028 0188 0001 0001 0021 ...... 00002 . . . 023 0004 0004 . . . 00029 1 8 1 2 000006 000012 0 000056 0 ...... 0 16 41 0 4 6 6 5 5 0 0 . 111000 0 2 0 . 0 27 0 0 . 10 0 0 0 . 0 0 . . . 342 0 0 0 10 2500 1117 0 0 0 0 . 0 × 0 0 0 0 +3 − 3180 0 ± ± ± ± 0 0 0 ± ± ± +27 − ± ± ± ± ± ± ± ± ± ± ± ± 0 ± ± ± ± ± ± ± ± ± ± ± ± +1999 − +1118 − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 38 ± Period Eccentricity . 1146 1951 4 3138 298.5 777.0 678.0 990.7 714.3 410.2 497.5 2073.6 1075.0 2354.3 1748.2 145.4 28.44 840.8 4451.8 528.07 379.63 801.30 396.03 446.27 466.47 1474.9 326.03 505.65 31.66 2136.14 621.99 ∼ 511.098 147.968 5.8953 9.0090 7929.40 11593.2 83.9152 90.2695 52.8657 62.6277 79.4179 9017.76 5.69449 14.18643 103.698 2.430420 1.328300 8.428198 55 5 09 97 2 6 1 3 7 1 9 5 0 7 28 0 56 6 22 0 8 0 4 59 93 6 0 1 5 02 . . 73 08 04 3 . . . 9 ...... 8 . 9 ...... i ) (day) 8 9 0 ...... 2 0 0 1 0 4 2 4 8 1 1 3 1 0 3 0 1 0 3 1 0 3 3 0 0 2 0 2 1 1 . 12 1 10 1 11 0 6 0 1 2 0 8 8 . . . sin ± ± ± ± ± ± ± ± ± ± ± ± ± +7 − ± ± ± ± ± +0 − ± +2 − ± ± ± ± ± ± ± ± ± ± ± Jup ± ± ± ± ± +20 − 6 6 3 ··· ··· ··· ··· ··· ··· ··· 77.8 15.0778 20.0 471.6 17.5 256.20 c 2 6 4 1 2 7 3 0 1 6 4 5 3 0 0 0 6 0 6 4 43 22 67 . . . Table 1. Close brown dwarfs (candidate) companions to solar 0 7 1 4 ...... M 23 99 65 28 13 37 99 64 17 . . . . 90 04 . . M ...... 53 25 59 55.1 . 19 62 59 46 13 22 27 28 68 50 57 64 40 60 47 37 25 30 19 20 12 39 14 10 13 14 13 18 13 10 38 15 0 24 9 8 2 . 0 1 . . . 8 . )( . . 2 0 5 . 15 . 9 . 0 19 13 14 26 . 1 24 c 5 ± ± +1 − +23 − ± Jup ± ± ± ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· M +26 − 3 8 2 4 9 . . M . . . ( 38 23 34 69 49 40 101 127b − Object HIP 67526 HD 30501 TYC 1240-945-1 HD 114762 HD 22781 HD137510 20.0-60.0 TYC 2930-872-1 HD 5388b HD 43848 HD 211847 HD 89707 HD 92320 TYC 3130-160 HIP 5158 HD 4747 HD 10697 BD+202457c TYC 2949-557-1 BD+202457b HD 140913 HD 189310 HD 180314 TYC 2087-255-1 HD 52756 HIP 78530 HD 65430 TYC 2534-698-1 HD 167665 HAT-P-13c HD 30246 HD 30339 HD 16760b NGC 4349 HD 41004B HD 13189 NGC 2423-3b HIP 21832 HD38529c 17.6 HD174457 107.8 HD 137759 GJ 595 HD 190228 HD 14651 HD 131664 HD 162020 HD 202206 HD 119445 HD 180777 HD 91669 11Com HD 191760 HD 168443

c 0000 RAS, MNRAS 000, 000–000 Statistical Properties of Brown Dwarf Companions: Implications for Different Formation Mechanisms 9 2007); 52 h (2000); 43 50 49 61 59, 60 13 15 14 19 20 23 56 iya et al. (2009); 9 ane et al. (2011); 35 ; 17 Halbwachs et al. ischer et al. (2002); 60 10 10 06 080 10 . . . . . 07 08 20 073 05 . . . . 043 0 0 0 0 0 . al. (2007); 26 Perrier et al. . 0 0 0 0 0 0 ± ± ± ± ± ± ± ± ± ± ± [Fe/H] References 14 03 02 29 . . . . 05 04 10 177 . . . 0 0 0 0 . 008 -0.01 . 0 -0.135 − − − − 0 − 2 0 2 10 06 0 09 01 11 07 008 0 041 019 ) 012 ...... g 0 0 0 0 0 0 0 0 0 0 0 ± ± +0 − ± ± ± ± ± ± ± ± log( 4.3 4.1 2.59 4.42 4.16 4.28 4.35 4.22 4.229 4.851 4.244 59 44 97 65 50 100 200 100 140 140 ± ± ± ± ± eff ± ± ± ± ± ··· 5060 4.57 -0.17 57 5835 4.64 -0.08 57 T 5883 6097 6201 5966 6518 ; 49 Creppet al. (2011); 50 Wang etal. (2012); 51 Patel etal. ( 4844 6350 6429 6740 6260 ? gt et al. (2000); 41 Wittenmyer et al. (2009); 42 Zucker & Maze l. (2010); 6 Wittenmyer et al. (2009); 7 Liu et al. (2008); 8 Om hy et al. (2011); 15 Johnson et al. (2011); 1610); Lee 32 et al. Niedzielski et (2011) al. (2009); 33 Frinket al. (2002); 34 K 0); 48 3 04 03 009 12 066 09 026 ...... 010); 57 Santerne et al. (2012); 58 Fleming et al. (2012); 59 F 07 03 ) (K) (cgs) 06 07 . . Bouchy et al. (2011); 24 Galland et al. (2006); 25 Gerbaldi et 0 0 0 0 0 0 0 0 . . ⋆ 0 0 ⊙ 0 ± ± ± ± ± ± ± ± ± ± M +0 − M 7 ( . 08 17 32 37 1.1 . . . . 1.24 1.02 370 213 324 1 . 1 1 . 1 . 1 01 05 008 2 032 0 023 022 0069 01 017 1 ...... 0 0 0 0 0 0 0 0.5 +0 − ± ± +0 − ± ± ··· ··· ··· ··· ± 0 (adopted) 0 (adopted) 0 (adopted) Table 1 (cont’d) 0.34 0.40 0.26 0.121 0.378 0.056 0.0099 2 . 5 7 0 . . . 0002 00004 4 00003 5 2 000013 000015 . 0000068 . 000005 . . . . 0 . ± 0 ± ± 0 1059 0 0 0 0 ± ± +803 − ± 3.66 ± ± ± 11.72 ± Period Eccentricity 1319 (2009); 38 Sato et al. (2009); 39 Mugrauer et al. (2006); 40 Vo 3 Hatzes et al. (2005); 4 Nidever et al. (2002); 5 Benedict et a 1366.8 1188.0 2010); 46 et Lafreni`ere al. (2011); 47 Konopacky et al. (201 26772 Ferozet al. (2011); 30 Zucker et al. (2004); 31 Winn et al. (20 21.0874 al. (2006); 55 Siverd et al. (2012); 56 Sozzetti & Desidera (2 006); 12 Udry et al. (2002); 13 Anderson et al. (2011); 14 Bouc 3.06036 014); 63 Jiang et al. (2013) 12.71382 4.25680 4.156736 1.217514 nn et al. (2008); 21 Ma et al. (2013); 22 Marcy et al. (2001); 23 3.1915239 8 . i ) (day) 8 47 . 2 . 0 0 sin ± ± Jup ± ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· c 3 . M M 13.0 12.0 37 0 89 3 1 00 59 . . . . . )( 48 5 . 91 93 3 0 2 4 1 . . . 0 . 0 18 c 0 ± ± ± ± ± +0 − +11 − Jup ± +0 − ··· ··· M 7 9 3 8 . . . M 42 96 66 . ( . . 68 62 63 11 60 21 Object HD 175679 HD 136118 HR 7672b References. — 1 Sahlmann et al. (2011a); 2 Sato et al. (2010); HD 217786 WASP-30 LHS 6343 COROT-15 XO-3b Corot-3b KOI-423b 18.0 KELT-1b 27.23 KOI-205.01 35 KOI-554.01 80 ¨ Lo Curto et al. (2010); 53 Santos et al. (2010); 54 Sozzetti et Moutou et al. (2009); 10(2000); Endl 18 et et D´ıaz al. al. (2012); (2004); 19 11(2003); Deleuil Butler et 27 et al. al. (2008); Jenkins 20 et (2 Wi al. (2009); 28 Correia et al. (2005); 29 Onehag et al. (2009); 36Latham et Lovis al. & Mayor (1989); (2007); 44 37 Kane et Santos al. et al. (2009); 45 Fleming et al. ( Martioli et al. (2010); 61 Moutou et al. (2011); 62 Ma et al. (2

c 0000 RAS, MNRAS 000, 000–000