Connecting substellar and stellar formation. The role of the host star's metallicity
AAstronomy & Astrophysics manuscript no. metal_planet_full c (cid:13)
ESO 2019March 5, 2019
Connecting substellar and stellar formation. The role of the hoststar’s metallicity.
J. Maldonado , E. Villaver , C. Eiroa , and G. Micela INAF - Osservatorio Astronomico di Palermo, Piazza del Parlamento 1, 90134 Palermo, Italye-mail: [email protected] Universidad Autónoma de Madrid, Dpto. Física Teórica, Facultad de Ciencias, Campus de Cantoblanco, 28049 Madrid, SpainReceived September 15, 1996; accepted March 16, 1997
ABSTRACT
Context.
Most of our current understanding of the planet formation mechanism is based on the planet metallicity correlation derivedmostly from solar-type stars harbouring gas-giant planets.
Aims.
To achieve a far more reaching grasp on the substellar formation process we aim to analyse in terms of their metallicity adiverse sample of stars (in terms of mass and spectral type) covering the whole range of possible outcomes of the planet formationprocess (from planetesimals to brown dwarfs and low-mass binaries).
Methods.
Our methodology is based on the use of high-precision stellar parameters derived by our own group in previous works fromhigh-resolution spectra by using the iron ionisation and equilibrium conditions. All values are derived in an homogeneous way, exceptfor the M dwarfs where a methodology based on the use of pseudo equivalent widths of spectral features was used.
Results.
Our results show that as the mass of the substellar companion increases the metallicity of the host star tendency is to lowervalues. The same trend is maintained when analysing stars with low-mass stellar companions and a tendency towards a wide rangeof host star’s metallicity is found for systems with low mass planets. We also confirm that more massive planets tend to orbit aroundmore massive stars.
Conclusions.
The core-accretion formation mechanism for planet formation achieves its maximum e ffi ciency for planets with massesin the range 0.2 and 2 M Jup . Substellar objects with higher masses have higher probabilities of being formed as stars. Low-massplanets and planetesimals might be formed by core-accretion even around low-metallicity stars.
Key words. techniques: spectroscopic - stars: abundances -stars: late-type -stars: planetary systems
1. Introduction
Exoplanetary science has succeeded in discovering an astonish-ing diversity of planetary systems. The role of the host star’smetallicity in planet formation has been largely discussed withthe finding that the frequency of giant planets is a strong functionof the stellar metallicity (Gonzalez 1997; Santos et al. 2004; Fis-cher & Valenti 2005). The planet-metallicity correlation is usu-ally interpreted in the framework of the core-accretion model(e.g. Pollack et al. 1996) as the final mass of cores via oligarchicgrowth increases with the solid density in proto-planetary discs(Kokubo & Ida 2002).Initially found for gas-giant planets around solar-type stars(Gonzalez 1997; Fischer & Valenti 2005), many works have triedto probe whether the gas-giant planet metallicity correlation alsoholds for other kind of stars as well as other types of substel-lar objects, e.g. low-mass planets (Ghezzi et al. 2010b; Mayoret al. 2011; Sousa et al. 2011; Buchhave et al. 2012; Buchhave& Latham 2015), brown dwarfs (Sahlmann et al. 2011; Ma & Ge2014; Mata Sánchez et al. 2014; Maldonado & Villaver 2017),stars with debris discs (Beichman et al. 2005; Chavero et al.2006; Greaves et al. 2006; Bryden et al. 2009; Kóspál et al. 2009;Maldonado et al. 2012, 2015b; Gáspár et al. 2016), evolved (sub-giant and red giant) stars (Sadakane et al. 2005; Schuler et al.2005; Hekker & Meléndez 2007; Pasquini et al. 2007; Takedaet al. 2008; Ghezzi et al. 2010a; Maldonado et al. 2013; Mortieret al. 2013; Jofré et al. 2015; Re ff ert et al. 2015; Maldonado & Villaver 2016), low-mass (M dwarf) stars (e.g. Neves et al.2013), or if there are di ff erences in the host star metallicitywhen close-in and more distant planets are present in the sys-tem (Sozzetti 2004; Maldonado et al. 2018; Wilson et al. 2018).It should be noted that most of these references refer to radialvelocity planets.As clear from the references above, previous works focuson particular types of stars and planets and, to the very best ofour knowledge, a global view of the planet-metallicity correla-tion and its implications on the planet formation process are stillmissing. This is precisely the goal of this work in which we anal-yse in the most homogeneous possible way a large sample ofstars harbouring the full range of possible outcomes of the planetformation process (from debris discs to massive brown dwarfs)and without any restriction of the host star’s spectral type (fromM dwarfs to early-F) or evolutionary status (from main-sequenceto giants). The analysis is completed with literature data of low-mass binary stars in order to set the results into a general context.The paper is organised as follows. Section 2 describes thestellar sample. The completeness of the planet host subsample isanalysed in Sect. 3. The analysis of the host star metallicities asa function of the substellar companion mass and the mass of thehost star is performed in Sect. 4. The results are discussed in thecontext of current planet formation models in Sect. 5. A com-parison with the results from the K epler mission is performed inSect. 6. Our conclusions follow in Sect. 7. Article number, page 1 of 18 a r X i v : . [ a s t r o - ph . S R ] M a r & A proofs: manuscript no. metal_planet_full
Table 1.
Architecture of the planetary systems in our stellar sample.
Type Number NotesSubstellar objects 345 95 multiple systems(total)Brown dwarfs 59 3 systems with 2 BDs( 10 M
Jup < m C sin i <
70 M
Jup ) 5 systems BD + planetLow-mass planets 78 34 hot (a < C sin i <
30 M ⊕ ) 44 cool (a > < C sin i >
30 M ⊕ ) 174 cool (a > + substellar object
2. Stellar sample
Our stellar sample is selected from our previous works (Mal-donado et al. 2012, 2013, 2015a,b, 2018; Maldonado & Villaver2016, 2017) which might be consulted for further details. Briefly,high-resolution échelle spectra of the stars were obtained in 2-3 metre class telescopes or obtained from public archives. Ba-sic stellar parameters (T e ff , log g , microturbulent velocity ξ t , and[Fe / H]) were determined by using the code TGVIT (Takedaet al. 2005) which implements the iron ionisation, match of thecurve of growth and iron equilibrium conditions. Stellar age,mass, and radius were computed from H ipparcos V magnitudes(ESA 1997) and parallaxes (van Leeuwen 2007), when avail-able, using the code PARAM (da Silva et al. 2006), togetherwith the PARSEC set of isochrones (Bressan et al. 2012). Forsome planet hosts where the PARAM code failed to give a rea-sonable value we took the mass values from the NASA exoplanetarchive (specifically from the summary of stellar information ta-ble, for all stars but KOI 415 for which we took the value fromthe KOI stellar properites table). For the M dwarfs a di ff erentmethodology was used. The analysis is based on the use of ra-tios of pseudo equivalent widths of spectral features which aresensitive to the e ff ective temperature and the stellar metallicity(Maldonado et al. 2015a). Stellar masses and luminosities for Mdwarfs were obtained from the derived temperatures and metal-licities by using empirical calibrations.The total number of stars in our sample amounts to 551. It iscomposed of 71 F-type stars, 261 G-type stars, 166 K-type stars,and 53 early-Ms. Regarding their evolutionary state, 373 are inthe main-sequence, 63 are classified as subgiants, and 115 aregiants.The di ff erent architectures of the planetary systems har-boured by our sample (type and number of substellar objects) areshown in Table 1, while figure 1 shows the Hertzsprung-Russell(HR) diagram of the stars analysed. Our sample is mainly com-posed of “mature” planetary systems as only ∼
7% of our starshave estimated ages younger than 500 Myr.
3. Sample completeness
A deep analysis of the completeness or detectability of our planethost sample is di ffi cult to overcome as our targets are not selectedfrom any particular exoplanet survey. Indeed, they were ini-tially selected to address di ff erent individual aspects of the planetformation process (presence of discs, planet formation aroundevolved stars, brown-dwarf vs. planet formation, hot vs. coolplanets ...). Our planet host sample includes targets from dif-ferent radial velocity surveys, e.g., the HARPS search for south-ern extra-solar planets (Pepe et al. 2004); the Anglo-Australianplanet search (Tinney et al. 2001); the N2K survey (Fischer http: // optik2.mtk.nao.ac.jp / ~takeda / tgv / http: // stev.oapd.inaf.it / cgi-bin / param Fig. 1.
Luminosity versus T e ff diagram for the stars analysed. Stars withdiscs are plotted as green whilst stars without discs are shown in blue.Stars with planets are shown as filled circles and stars with companionsin the brown dwarf regime are shown in filled stars. Some evolutionarytracks ranging from 0.7 to 3.0 solar masses from Girardi et al. (2000)are overplotted. For each mass, three tracks are plotted, correspondingto Z = / H] = -0.4 dex, dotted lines), Z = / H] = + = / H] = + et al. 2005); the UCO / Lick survey (Hekker et al. 2006); the thePennState-Toru´n Centre for Astronomy Planet Search (Niedziel-ski & Wolszczan 2008); the retired A stars project (Johnson et al.2007); or the list of stars with brown dwarf companions by Ma& Ge (2014) and Wilson et al. (2016); among others. In otherwords, our planet host sample comes from a wide variety ofplanet search programmes with (most likely) di ff erent selectioncriteria, sensibilities, and biases, sampling significantly di ff erentregions of the HR diagram.On the other hand, the comparison sample is mainly drawnfrom the H ipparcos catalogue (ESA 1997) and was chosen tocover similar stellar parameters as the stars with detected planets.For the sake of completeness, we give in Table A.1 the basicproperties of the full sample of stars covered here. Further detailscan be found in our previous works (see references above).In order to estimate the detectability limits of our planet hostswe proceeded as follows. For each star we searched for its cor-responding radial velocity curve. Whenever possible, radial ve-locity data was taken from the NASA Exoplanet Archive . Oth-erwise, we searched for the data in the corresponding discov-ery’s paper. We were able to recover the radial velocity seriesfor 89.6% of our stars with planets. We subtracted the contri-bution of the known planets to each radial velocity data set byfitting a keplerian orbit using the code rvlin (Wright & Howard2009). The fits were done by fixing the planetary period to thepublished values. When several planets were present around thesame star, we subtract them in a sequential way. We took intoaccount the fact that data obtained with di ff erent instrumentsmight be available for the same star. Once the contribution forthe known planets have been removed from the radial velocity https: // exoplanetarchive.ipac.caltech.edu http: // exoplanets.org / code / Article number, page 2 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity. datasets, we considered the rms of the residuals to be represen-tative of our measurement uncertainty (e.g. Endl et al. 2001).For each planet host we computed the expected radial ve-locity semiamplitude due to the presence of di ff erent types ofplanets considering circular orbits, as usually done in the liter-ature (see e.g. Mayor et al. 2011, and references therein). Wealso note that it has been shown that even eccentricities as highas 0.5 do not have a strong influence in the planet detection’slimits (Endl et al. 2002; Cumming & Dragomir 2010). We sam-pled the planetary mass space in logarithmic space, with valuesranging from 0.005 to 80 M Jup . Regarding the orbital periods,we sampled the orbital frequencies (also in logarithmic space)from periods from one to 10 days. For each planet we com-puted the expected radial velocities keeping the same time as theoriginal observations. A total of eleven realizations of the radialvelocity, each simulation corresponding to a di ff erent phase o ff -set (from 0 to 2 π ), were performed. We considered a planet tobe detectable around a given star if the rms of the planet’s ex-pected radial velocity is larger than the rms of the stellar radialvelocity residuals in each of the simulated phases (Galland et al.2005; Lagrange et al. 2012; Meunier et al. 2012). We are awarethat this is a “conservative” approach, i.e, it might overestimatethe detection limits for some periods. It is however, a fast androbust method, ideal for achieving a quick look an for obtainingan e ffi cient determination of the detection limits (Meunier et al.2012). It should be noted that for the scope of this work a con-servative approach should be preferred in order to obtain robustconclusions.Figure 2 shows the derived detection probability curves.They show for each period, the percentage of stars from oursample for which planets with the corresponding minimum massmight be detected, i.e., planets located in the region above the p % curve can be detected in p % of our stars. Fig. 2.
Minimum mass versus planetary period diagram. Substellarcompanions analysed in this work are shown in grey circles. Detectionprobability curves are superimposed with di ff erent colours. Horizontaldashed lines indicate the standard mass loci of low-mass planets, andgas-giant planets companions. Several main conclusions can be drawn from this plot: i)First, most of the long-period planets ( P >
100 days) are de-tectable in approximately more than 60-80% of our targets. Onlyfor planets with periods longer than ∼ ∼
40% and below; ii) Low-mass planets, on theother hand, are detectable only in a small fraction of stars, be-tween 2 and 20%; iii) Finally, it can be seen that planets with themass of Jupiter and short periods are detectable in practically allstars.As noted before our approach is quite conservative, so it isnot surprising that some planets are actually located in the regionunder the 2% probability curve. We will discuss at length theimplications of these findings in our analysis in the next sections.Figure 3 shows the detectability limits for di ff erent subsam-ples of interest, see Sect. 4. In addition to the conclusions fromthe previous figure it can be seen the detectability curves for sur-veys aiming to detect low-mass planets are clearly shifted to-wards lower planetary mass companions, as expected. On theother hand surveys of stars with brown dwarfs and specially sur-veys of giant stars show significant higher detection limits. Starswith cool and hot Jupiters show nearly identical detection curves.Finally, we can see that we are mostly insensitive to the presenceof small planets (specially evident in the case of surveys aroundevolved stars) which demonstrates the need of dedicated inten-sive long-term surveys (and probably the development of spe-cific techniques to deal with the stellar noise problem) in orderto detect this kind of planets.
4. Analysis
Figure 4 compares the cumulative distribution function of thestellar metallicity of the di ff erent stars analysed in this work. Thestars have been divided into stars hosting hot Jupiters (if planetsare located at distances smaller than 0.1 au), stars hosting cooldistant-gas-giant planets, giant stars with planets, stars harbour-ing low-mass brown dwarfs (m C sin i < Jup ), stars withhigh-mass brown dwarfs (m C sin i > Jup ), stars harbour-ing only low-mass planets (m C sin i <
30 M ⊕ ), and stars host-ing only debris discs. The figure shows that while the metallic-ity distribution of stars with hot and cool gas-giant planets areshifted towards high metallicity values, this is not the case forthe other samples, which show metallicity distributions consis-tent with that of the comparison sample (i.e., stars without sub-stellar companions).Figure 5 shows the host star metallicity as a function of the(minimum) mass of the substellar companion. Colours and sym-bols indicate the mass of the host star. The figure clearly showsa tendency of lower host star’s metallicities as the mass of thesubstellar companion increases. The figure also shows that moremassive planets tend to orbit around more massive stars. It isclear from the figure that there is lack of planets around starswith metallicities lower than ∼ -0.4 dex. Stars with metallicitiesbelow this limit only harbour substellar companions in the browndwarf regime. We note that this result reefers to our sample andseveral planetary companions around more metal poor stars havebeen found.Our results suggest that there is a non universal planet forma-tion mechanism. Di ff erent mechanisms may operate altogetherand their relative e ffi ciency change with the mass of the substel-lar object that is formed. For substellar objects with masses inthe range 30 M ⊕ - 1 M Jup , high host star metallicities are found,suggesting that these planets are mainly formed by the core-accretion mechanism. As we move towards more massive sub-
Article number, page 3 of 18 & A proofs: manuscript no. metal_planet_full
Fig. 3.
Same as Fig. 2 but for the di ff erent subsamples described in Sect. 4. Fig. 4. [Fe / H] cumulative frequencies for the di ff erent samples analysedin this work. stellar objects, the range of the host star metallicities increasestowards more negative values, suggesting that a non-metallicity dependent formation mechanisms, such as gravitational instabil-ity or gravoturbulent fragmentation, might be at work.In a recent work, Schlaufman (2018) computed the massat which substellar companions no longer preferentially orbitmetal-rich stars finding that while objects with masses below 10M Jup orbit metal-rich stars, substellar companions with masseslarger than 10 M
Jup do not orbit metal-rich stars. We believethat our results are compatible with the findings by Schlaufman(2018) showing that the most massive substellar objects tend toform like stars.Figure 6 shows the orbital period of the substellar compan-ions as a function of the (minimum) mass of the substellar com-panions. Di ff erent colours indicate the eccentricity. The figureshows that the more massive substellar companions show largerperiods and eccentricities (P >
100 days, e > i) theirhost stars show a wider (towards negative values) range of metal-licities and higher stellar masses; ii) planets (or brown dwarfs)show longer periods and higher eccentricities. The di ff erences inperiod and eccentricity distributions between both types of plan-ets might be indicative of a di ff erent formation mechanism. Inaddition, the trend with the host star metallicity suggests that thehigher the mass of the substellar companion, the higher the prob-ability that it is formed by a non-metallicity dependent formation Article number, page 4 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Fig. 5.
Stellar metallicity of the host stars as a function of the minimum mass of the substellar companions. Di ff erent colours and symbol sizesindicate the mass of the host star. Vertical dashed lines indicate the standard mass loci of low-mass planets, gas-giant planets, brown dwarf, andstellar companions, from left to right respectively. Fig. 6.
Orbital period as a function of the minimum mass. Di ff er-ent colours and symbol sizes indicate the eccentricity values. Verticaldashed lines indicate the standard mass loci of low-mass planets, gas-giant planets, and brown dwarfs, from left to right. mechanism. This general trend explain many of the correlationsbetween the host star’s metallicity and planetary properties dis-cussed in recent works. – More massive stars host more massive planets.
It has beennoticed that giant stars host more massive planets than theirmain-sequence counterparts (e.g. Johnson et al. 2007; Lovis& Mayor 2007; Maldonado et al. 2013), although this resultshould be taken with caution as the detection of small plan-ets around evolved stars is hampered by the large levels ofstellar jitter in these stars (e.g. Niedzielski et al. 2016). Wenote that in our sample, stars in the mass range 1.5-2 M (cid:12) host only planets with masses around 1 M
Jup , while for starsmore massive than 2 M (cid:12) planets are more massive than 2M
Jup . This trend might reflect a correlation between disc gasmasses and giant planet masses (Alibert et al. 2011; Mor-dasini et al. 2012) as high-mass stars are likely to harbourmore massive protoplanetary disk (e.g. Natta et al. 2000). Inthis scenario giant planet formation can occur in low metal-licity but high-mass protoplanetary discs as it is the amountof metals in the disc the factor that drives the planet forma-tion process (e.g. Ghezzi et al. 2018). The metallicity e ff ectwould depend on the mass of the disc, being the minimummetallicity required to for a massive planet lower for massivestars. Ghezzi et al. (2018) found that the relation between the amount of metals in the protoplanetary disc and the forma-tion of giant planets does almost follow a linear relationship.The lack of a clear planet-metallicity correlation found forgiant stars might be explained by the fact that they host moremassive planets and these planets might find a way in theirmore massive planetary discs to bypass the core-acretionmechanism and form more like stars. Finally, we note a ten-dency of more massive giants stars with substellar compan-ions to have higher metallicities in agreement with previousworks (Maldonado et al. 2013; Jofré et al. 2015). – Trends in brown dwarfs hosts.
Ma & Ge (2014); MataSánchez et al. (2014) showed that unlike gas-giant planethosts, stars with brown dwarfs do not show metal-enrichment. Maldonado & Villaver (2017) found that starswith low-mass brown dwarfs tend to show higher metallici-ties than stars hosting more massive brown dwarfs. Ma & Ge(2014); Maldonado & Villaver (2017) also discussed di ff er-ences in the period-eccentricity distribution of massive andlow-mass brown dwarfs. This result fits well with our inter-pretation that more massive substellar objects tend to formmore like stars. – Close-in and more distant planets.
Recent works (Sozzetti2004; Maldonado et al. 2018; Wilson et al. 2018) have dis-cussed whether hot Jupiters host stars show higher metallic-ities than more distant planets. As more distant planets aremore massive than hot Jupiters (Ribas & Miralda-Escudé2007; Bashi et al. 2017; Santos et al. 2017; Jenkins et al.2017; Maldonado et al. 2018), see also Figure 6, they tend toorbit stars with a wider range of metallicities. – Planets around low-mass stars.
Planets around low-massstars (M (cid:63) < (cid:12) ) are mainly low-mass planets and theirhost stars do not show metal enrichment. They have shortperiods and low eccentricities. On the other hand, very fewgas-giant planets have been found orbiting around low-massstars, showing their host stars metal-enrichment (Neves et al.2013). We caution that these results refer to radial velocityplanets. Results from transit surveys are discussed in Sect. 6.We do not expect the general metallicity trends discussed inthis work for massive planets and brown dwarfs to be severelya ff ected by the di ff erent detection limits achieved for the di ff er-ent planet hosts. As discussed in Sect. 3 planets of the mass ofJupiter at short periods can be detected in more than 95% of Article number, page 5 of 18 & A proofs: manuscript no. metal_planet_full our targets. More distant substellar companions ( P >
100 days)might be detected in a significant large percentage of our stars,between 60 and 80%. The possible trend between massive plan-etary companions and large periods might however be a ff ectedby our lower sensitivity to detect small gaseous planets (withmasses of the order of the mass of Jupiter) at very long periods( P >
5. The planet-metallicity correlation in context
In order to discuss our results into a broader context, data fromlow-mass binaries have been included in Figure 5. The data istaken from Mann et al. (2013) who compiled the metallicity ofthe primary stars mainly from high-resolution spectra. The massof the late-K, M companions as well as the primaries are esti-mated by using a spectral type - stellar mass relationship basedon the data provided by Cox (2000).The figure shows that the tendency of a wider range of metal-licities (lower values) towards more massive objects continuesin the low-mass stellar range. Despite the fact we are comparingthe minimum mass of substellar objects with estimates of themass of stars, the trend that more massive substellar or stellarobjects tend to form in a non-metallicity dependent mechanismseems to hold, i.e., there seems to be a continuity between sub-stellar and stellar companions. According to this figure the core-accretion mechanism for planet formation would have its high-est e ffi ciency for forming planets with masses around 1 Jupiter’smass (hot Jupiters).It has been shown that the fraction of close binaries of solar-type stars decreases with the metallicity while the wide binaryfraction is basically constant with metallicity at large separa-tions (e.g. Moe et al. 2018; El-Badry & Rix 2018). Followingthe reasoning of Moe et al. (2018) massive and close substel-lar companions might form by fragmentation of the protostellardisc. Protostellar discs of solar-type stars are usually opticallythick and lower metallicities imply lower opacities and enhancedcooling rates which translate in higher probabilities of disc frag-mentation . On the other hand, massive and distant companionsmight form by turbulent fragmentation of molecular cores, a pro-cess which is known to be independent of metallicity.Figure 5 also reveals a possible tendency of wider metal-licities towards low-mass planets which may still be formed bycore-accretion around low-metallicity stars. The low-metallicityenvironment implies long times for forming a core able to accretgas before the disc’s dissipation, so only small planets and plan-etesimals can be formed (e.g. Mordasini et al. 2012). However,the sample of M dwarfs planet hosts is still too small to makea strong claim in these sense. We also should note our limitedsensitivity to low-mass planet’s detection. Debris disc’s massesdo not help either as they are usually unavailable and subject tomany assumptions.
6. Comparison with Kepler results
Given that our planet host sample is mainly selected from radialvelocity surveys a comparison with the results from the K epler Although for very low metallicities, the disc becomes optically thinand the e ff ect of lower metallicity would be the opposite. mission is mandatory in order to achieve a full vision of planetformation.In a recent work, Petigura et al. (2018) analyse a largesample of K epler objects of interest with metallicities derivedfrom spectroscopic observations finding that planets smaller thanNeptunes (R P < ⊕ ) are found around stars with a wide rangeof metallicities. On the other hand, sub-Saturns and Jupiters arefound around metal-rich stars (their Figure 3). The authors alsonote a gradual upward trend in mean host star metallicity fromsmaller to larger planets in agreement with previous analysis ofsmaller samples (Buchhave et al. 2012, 2014). These results sup-port our findings that only stars hosting Jupiter-like planets showpreferentially the metal-rich signature. As lower planetary ra-dius implies lower planetary masses, although the relationship iscomplex and depend on the planet composition (e.g. Lopez &Fortney 2014), we conclude that the K epler data supports oursuggestion that a tendency of lower metallicities towards low-mass planets might be hidden in Fig. 5 as discussed in Sect. 5.Similar results have been found by Narang et al. (2018) whoshow that the host star metallicity, increases with larger planetaryradius / mass up to about 1 M Jup or 4 R ⊕ . For planetary masseslarger than 4 M Jup the authors also found that more massive plan-ets have on average lower host star metallicities in agreementwith our findings. The authors also discuss that hot transitingplanets (periods less than 10 days) orbit around stars with higheraverage metallicity in agreement with our previous results (Mal-donado et al. 2018) and this work.Studies of the K epler occurrence rates (Mulders et al.2015a,b) have confirmed that small planets (1.0 - 3.0 R ⊕ ) aremore common around M dwarfs than around main-sequenceFGK stars (Howard et al. 2012). At larger planetary radii planetsbecome more common around sun-like stars. Despite begin dif-ferent samples (1.0 - 3.0 R ⊕ planets correspond to masses below ∼ ⊕ , i.e., planets smaller than Neptune), a similar tendency ofa larger occurrence of small planets towards less massive stars isfound in our results from Fig. 5 where it can be seen that the vastmajority of the low-mass planets (m C sin i < ⊕ ) orbit aroundstars with masses below 1 M (cid:12) .
7. Conclusions
Achieving a full vision of how planets and planetary systemsform and evolve is only possible by analysing in a homogeneousway large samples of stars covering the full domain of param-eters, i.e, including the di ff erent outcomes of the planet forma-tion process (from planetesimals to massive brown dwarfs andlow-mass stars) as well as the full range of host star’s massesand types. In this work we performed a detailed analysis of theplanet-metallicity correlation by analysing in a joined way thedata from our previous works, focused on certain types of starsand / or planets. Most of the studied stars (excluding the M dwarfsubsample) was analysed in the same way using similar spectraand techniques.Our results show a continuity between the formation ofsubstellar and stellar companions driven by the metallicity ofthe host star. The core-accretion formation mechanism wouldachieve its maximum e ffi ciency for planets with masses between ∼ Jup . For more massive substellar objects as well inlow-mass binary companions the range of host star’s metallici-ties increases towards lower values, suggesting that both kind ofobjects tend to share similar formation mechanisms.Another tendency towards lower host star’s metallicitiesseems to be present towards the less massive outcomes of the
Article number, page 6 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity. planet formation process (low-mass planets and probably plan-etesimals) which may still be formed by the core-accretionmethod. However, this tendency might need additional confir-mation.
Acknowledgements.
This research was supported by the Italian Ministry of Ed-ucation, University, and Research through the
PREMIALE WOW 2013 researchproject under grant
Ricerca di pianeti intorno a stelle di piccola massa . J. M.acknowledges support from the Ariel ASI-INAF agreement N. 2015-038-R.0.E. V., and C. E. acknowledge support from the
On the rocks project funded bythe Spanish Ministerio de Economía y Competitividad under grant
AYA2014-55840-P . We sincerely appreciate the careful reading of the manuscript and theconstructive comments of an anonymous referee.
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Appendix A: Additional tables
Article number, page 8 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Table A.1. Basic properties of the stars considered in this work.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )171 224930 5.80 G3V -0.83 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / G3IV / V 0.29 ± ± ± ± ± ± ± ± ± ± / G5V -0.67 ± ± ± ± ± ± ± ± ± / K0V -0.18 ± ± ± ± ± ± ± ± ± ± ± ± + planet5336 6582 5.17 G5V -0.87 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ‡ ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 9 of 18 & A proofs: manuscript no. metal_planet_full
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )12114 16160 5.79 K3V -0.10 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / K0V -0.68 ± ± ± ± ± ± ± ± ± ± ± ± ‡ ± ± ± ± ± ± ‡ ± ± / G0V -0.16 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / V -0.14 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / K0V 0.17 ± ± ± ± ± ± / G5V 0.27 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 10 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )28767 40979 6.74 F8 0.18 ± ± ± ± ± ± ± ± + low-mass planet29568 43162 6.37 G5V 0.01 ± ± ± ± ± ± ± ± ± ± / G2V -0.17 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / K2III 0.00 ± ± a) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± b) / K0V 0.26 ± ± ± ± ± ± / G5V 0.26 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 11 of 18 & A proofs: manuscript no. metal_planet_full
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )44291 77338 8.63 K0IV 0.38 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± + ± ± c) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / G5V -0.06 ± ± ± ± ± ± ‡ ± ± ± ± / G8V -0.86 ± ± ± ± ± ± ± ± ± ± ± / G5V -0.31 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 12 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )58952 104985 5.79 G9III -0.36 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ‡ ± ± ± ± / G0V 0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / V 0.22 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 13 of 18 & A proofs: manuscript no. metal_planet_full
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )70695 126525 7.85 G5V -0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / G5V 0.03 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 14 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )81819 150474 7.16 G8V -0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ‡ ± ‡ /
6? low-mass planets84856 156846 6.50 G1V 0.13 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / IV 0.01 ± ± Article number, page 15 of 18 & A proofs: manuscript no. metal_planet_full
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )95124 181342 7.55 K0III 0.20 ± ± / G2V -0.11 ± ± ± ± ± ± ± ± + BD95822 183492 5.57 K0III 0.07 ± ± ± ± ± ± ± ± ± ± ‡ ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / V 0.24 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / K0V -0.25 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 16 of 18. Maldonado et al.: Connecting substellar and stellar formation. The role of the host star’s metallicity.
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )105312 202940 6.56 G5V -0.31 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / K1V -0.30 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 17 of 18 & A proofs: manuscript no. metal_planet_full
Table A.1.
Continued.
HIP / Other HD V SpType [Fe / H] M (cid:63)
Ref † Sample / Notes(mag) (dex) (M (cid:12) )116584 222107 3.81 G8III-IV -0.42 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Notes. † Reference for metallicity and stellar mass: (1)
Maldonado et al. (2015b); (2)
Maldonado & Villaver (2016); (3)
Maldonado & Villaver(2017); (4)
Maldonado et al. (2018); (5)
Maldonado et al. (2015a); ‡ Value from the NASA exoplanet archive. Specifically from the summaryof stellar information table for all stars, but KOI 415 for which we took the value from the KOI stellar properites table; a) Lovis & Mayor(2007); b) Bouchy et al. (2016); c)c)