Searching for the Signatures of Terrestial Planets in Solar Analogs

Searching for the signatures of terrestial planets in solar analogs

J. I. Gonz´alez Hern´andez1,2, G. Israelian1, N. C. Santos3,4, S. Sousa3, E. Delgado-Mena1, V. Neves3, & S. Udry5 [email protected]

ABSTRACT We present a fully differential chemical abundance analysis using very high- resolution (λ/δλ & 85, 000) and very high signal-to-noise (S/N ∼ 800 on average) HARPS and UVES spectra of 7 solar twins and 95 solar analogs, 24 are planet hosts and 71 are stars without detected planets. The whole sample of solar analogs provide very accurate Galactic chemical evolution trends in the metalliciy range −0.3 < [Fe/H] < 0.5. Solar twins with and without planets show similar mean abundance ratios. We have also analysed a sub-sample of 28 solar analogs, 14 planet hosts and 14 stars without known planets, with spectra at S/N ∼ 850 on average, in the metallicity range 0.14 < [Fe/H] < 0.36 and find the same abundance pattern for both samples of stars with and without planets. This result does not depend on either the planet mass, from 7 Earth masses to 17.4 Jupiter masses, or the orbital period of the planets, from 3 to 4300 days. In addition, we have derived the slope of the abundance ratios as a function of the condensation temperature for each star and again find similar distributions of the slopes for both stars with and without planets. In particular, the peaks of these two distributions are placed at a similar value but with opposite sign as that expected from a possible signature of terrestial planets. In particular, two of the planetary systems in this sample, containing each of them a Super-Earth like planet, show slope values very close to these peaks which may suggest that these abundance patterns are not related to the presence of terrestial planets.

Subject headings: stars: abundances — stars: fundamental parameters — stars: planetary systems — stars: planetary systems: formation — stars: atmospheres

1Instituto de Astrof´ısica de Canarias, C/ Via Atm´osfera, Facultad de Ciencias F´ısicas, Univer- L´actea s/n, 38200 La Laguna, Spain: [email protected] sidad Complutense de Madrid, E-28040 Madrid, 2Dpto. de Astrof´ısica y Ciencias de la Spain 3Centro de Astrof´ısica, Universidade do Porto,

1 1. Introduction differences in some element abundance ratios have been suggested (Bodaghee et al. The discovery of first exoplanet orbiting 2003; Beirão et al. 2005; Gilli et al. 2006; a solar-type star by Mayor & Queloz (1995) Robinson et al. 2006; Neves et al. 2009), initiated a new and very attractive field although no definite explanation and/or in which the number of studies is contin- conclusion has been established yet due uously increasing. A substantial amount to the contradictory results among these of spectroscopic data has been collected studies. See also the reviews on chemi- since then which have allowed to perform cal abundance trends in planet-host stars not only radial velocity planet searches (see by Israelian (2006, 2007, 2008) and Santos e.g Udry & Santos 2007; Udry & Mayor (2006, 2008). 2008) but also chemical abundance analy- Smith et al. (2001) firstly examined the sis (e.g Gonzalez et al. 2001; Sadakane et correlations between the abundance ratio al. 2002; Ecuvillon et al. 2004; Ecuvillon [X/H] as a function of the condensation et al. 2006a; Gilli et al. 2006; Neves et al. temperature, T , and they found only six 2009) and the study of kinematic proper- C stars among a sample of 20 planet-host ties (e.g. Santos et al. 2003; Ecuvillon et stars from Gonzalez et al. (2001) which al. 2007) trying to focus on finding differ- have positive correlations, suggesting pos- ent signatures able to distiguish stars with sible accretion of planetesimals. Later on, and without known planets. Ecuvillon et al. (2006b) showed the dis- The metal-rich nature of star hosting gi- tributions of the correlation [X/H] versus ant planets was firstly discussed by Gonza- TC in a sample of 88 planet-host stars lez (1997, 1998) and later on proved by a and 33 stars without known planets, but uniform analysis of large samples of planet- they did not find any significant differ- host stars and “single” (hereafter, “sin- ence between stars with and without plan- gle” refers to stars without known planets. However, the effective temperature ets) stars (Santos et al. 2001). Further range in this study was probably too large, studies has confirmed that the probabil- 4700 . Teff [K] . 6400, what may have ity of finding a planet strongly correlates smoothed out any possible existing signa- with the metal content of the parent star ture. Other studies, but on light elements, (e.g Santos et al. 2004, 2005; Valenti & have been able to find differences between Fisher 2005; Sousa et al. 2008). Following planet hosts and stars without planets (e.g. these studies and using the same spectro- Israelian et al. 2004; Takeda et al. 2007; scopic data, more complete chemical abun- Gonzalez et al. 2008). For instance, Is- dance studies have been done and small raelian et al. (2009) found that appar- ently Li is always very depleted in solar Rua das Estrelas, 4150-762 Porto, Portugal twins hosting planets whereas this is not 4Departamento de F´ısica e Astronomia, Facul- dade de Ciências, Universidade do Porto, Portugal the case for similar stars without detected 5Observatoire Astronomique de l’Université planets. This may suggest that planet for- de Genève, 51 Ch. des Maillettes, -Sauverny- mation may alter the Li content of a star. Ch1290, Versoix, Switzerland This result has been confirmed by Sousa et

2 Table 1 Spectroscopic observations

a b c d Spectrograph Telescope Spectral range λ/δλ Binning Nstars | S/N | δ | S/N | ∆S/N [A]˚ [A/pixel]˚

HARPS 3.6-m 3800–6900 110,000 0.010 85 886 450 370–2020 UVES VLT 4800–6800 85,000 0.015 8 829 173 550–1100 UES WHT 4150–7950 33,000 0.040 2 1400 – –

Note.—Details of the spectroscopic data used in this work. The UVES and UES stars are all planet-hosts. Two of the UVES stars were observed with a slit of 0.3′′ providing a resolving power λ/δλ ∼ 110, 000. aTotal number of stars including those with and without known planets. bMean signal-to-noise ratio, at λ = 6000 A,˚ of the spectroscopic data used in this work. cStandard deviation from the mean signal-to-noise ratio. dSignal-to-noise range of the spectroscopic data. al. (2010) and Gonzalez et al. (2010). the solar twins 16 Cyg A and 16 Gyg In the last decade, many studies have B, i.e. two stars with stellar parameters addressed the question whether the chemi- and metallicities very close to those of the cal abundances of the Sun are typical for a Sun. They used high-quality spectroscopic solar-type star, i.e. a star of solar mass and data of this binary star, and found very age (e.g. Guftafsson et al. 1998; Allende- small differences in their stellar parameters Prieto et al. 2006). The distribution of and Fe abundances. Later on, other au- spectroscopic metallicities of G-type stars thors have proposed and analyzed other so- in the solar neighbourhood has been esti- lar twins using a similar methodology (see mated to be ∼ −0.1 with a typical disper- Takeda 2005; Meléndez et al. 2006, 2007). sion of 0.2 dex (e.g. Edvarsson et al. 1993; More recently, Meléndez et al. (2009) have Allende-Prieto et al. 2004). These authors studied a sample of 11 solar twins. They also found small offsets in some element also used spectroscopic data at high re- abundances with respect to their solar val- solving power, λ/δλ = 65, 000, and high ues. However, these offsets are not always signal-to-noise ratio (S/N ∼ 450 per pixel). consistent among different studies which They observed some asteroids to use their brings some suspicion regarding their exis- spectra as a solar reference. By perform- tence and size. Allende-Prieto et al. (2006) ing a fully line-by-line differential analysis, analyzed a sample of solar analogs and did they obtained a clear trend, in the mean not find any significant offsets for Si, C, abundance differences of solar twins with Ca, Ti and Ni, suggesting that the offsets respect to the Sun, ∆[X/Fe]SUN−TWINS, as reported in previous studies were likely the a function of the condensation tempera- result of systematic errors. ture, TC , of the elements. The size of this Recently, Laws & Gonzalez (2001) per- trend is roughly 0.1 dex and goes from car- ∼ formed a differential abundance study of bon at ∆[C/Fe] 0.05 and TC = 40 K

3 to zirconium at ∆[Zr/Fe] ∼ −0.03 and provide the wavelength ranges, resolving TC = 1736 K. In spite of this tiny signa- power, S/N and other details of these spec- ture, they concluded that the most likely troscopic data. explanation to this abundance pattern is The data were reduced in a standard related to the presence of terrestial plan- manner, and latter normalized within ets in the solar planetary system. Ram´ırez IRAF1, using low-order polynomial fits to et al. (2009) confirmed this result on a the observed continuum. sample of 64 solar analogs, but with lower signal-to-noise spectra (S/N ∼ 200), al- 3. Sample though the correlation was not so clear. The discovery of more than 400 exoplan- To perform a detailed chemical analy- ets orbiting solar-type stars detected by sis, we selected stars with S/N > 350 from the radial velocity technique have provided these three spectrographs. We end up with a substantial sample of high-quality spec- a sample of 24 planet-host stars and 71 troscopic data, in particular, the HARPS “single” stars (see Fig. 1). All UVES and GTO planet search program which con- UES stars are planet hosts. These 95 solar tains so far about 450 stars (see e.g. Neves analogs of the sample have stellar param- et al. 2009). The multiplicity of these eters in the ranges 5600 < Teff [K] < 5954 −2 planetary systems is generating a certain and 4.0 < log(g[cm s ]) < 4.6, and metal- − number of questions about the processes licities in the range 0.3 < [Fe/H] < 0.5 of planet formation and evolution (see the (see Table 2). In Fig. 1, we display the introduction in Santos et al. 2010). histograms of the effective temperatures, surface gravities and metallicities of the In the present study, we will use the whole sample. The stars hosting planets high-quality spectroscopic data to get very are more metal-rich than the stars without accurate chemical abundances in a rela- known planets. Thus, we defined a new tively large sample of 95 solar analogs with sample of metal-rich solar analogs with and without planets, and to examine the 0.14 < [Fe/H] < 0.36, containing 14 planet results reported in Meléndez et al. (2009). hosts and 14 “single” stars. We have also 2. Observations considered the case of solar twins with and without planets. In that case, the number In this study we have made use of high- of stars goes down to only 2 planet-host resolution spectroscopic data obtained stars and 5 “single” stars. The ranges of with three different telescopes and instru- the stellar parameters and metallicities are ments: the 3.6-m telescope equipped with shown in Table 2.

HARPS at the Observatorio de La Silla in 1 Chile, the 8.2-m Kueyen VLT (UT2) tele- IRAF is distributed by the National Optical Ob- servatory, which is operated by the Association scope equipped with UVES at the Obser- of Universities for Research in Astronomy, Inc., vatorio Cerro Paranal in Chile and the 4.2- under contract with the National Science Founda- m WHT telescope equipped with UES at tion. the Observatorio del Roque de los Mucha- chos in La Palma, Spain. In Table 1, we

4 20 20 20 was computed using the 2002 version of the LTE code MOOG (Sneden 1973), and

15 15 15 a grid of Kurucz ATLAS9 plane-parallel model atmospheres (Kurucz 1993). For the HARPS spectra we just adopted the 10 10 10 published results in Sousa et al. (2008). For the UVES and UES stars, they had al- NUMBER OF STARS 5 5 5 ready spectroscopic stellar parameters reported in previous works (see Santos et al.

0 0 0 2004, 2005) computed from a set of 39 Fe I -0.3 0.0 0.3 5600 5800 6000 4.2 4.5 [Fe/H] TEFF log (g) and 12 Fe II and using FEROS, SARG and CORALIE spectra. Therefore, we decided Fig. 1.— Histograms of the stellar param- to re-compute the stellar parameters using eters and metallicities of the whole sample the higher quality data presented in this of solar analogs hosting planets (solid lines) study and the method described before. and without known planets (dashed lines). The new stellar parameters are consistent to those in Santos et al. (2004, 2005) al- Table 2: Ranges of stellar parameters and though the uncertainties are smaller due to metallicities in diﬀerent samples the higher quality of the new data reported in this work. In Table 3 we provide the new a b Sample Teﬀ log g [Fe/H] Nstars stellar parameters of the UVES and UES [K] [dex] [dex] stars. We note here that the similarities of SA 5600–5954 4.14–4.60 −0.30 – +0.50 95 mrSA 5600–5954 4.24–4.60 +0.14 – +0.36 28 the parameters of the stars in the sample ST 5700–5854 4.34–4.54 −0.07 – +0.07 7 with respect to the Sun allow us to achieve

a Stellar samples: solar analogs, “SA”, metal-rich solar a high accuracy in the chemical analysis analogs, “mrSA”, and solar twins “ST”. since the possible uncertainties on the at- b Total number of stars including those with and without known planets. mospheric model and on the stellar parameters and abundances are minimized. 4. Stellar Parameters 5. Chemical abundances The stellar parameters and metallicities The chemical analysis is done by com- of the whole sample of stars were computed puting the EWs of spectral lines using the using the method described in Sousa et code ARES (Sousa et al. 2007) for most al. (2008), based on the equivalent widths of the elements. We follow the rules given (EWs) of 263 Fe I and 36 Fe II lines, in Sousa et al. (2008) to adjust the ARES measured with the code ARES2 (Sousa et parameter rejt for each spectrum taking al. 2007) and evaluating the excitation into account the S/N ratio. We ﬁxed the and ionization equilibria. The chemical other ARES parameters to: smoother = 4, abundance derived for each spectral line space = 3, lineresol = 0.07, miniline = 2. 2 However, for O, S and Eu, we “manually” The ARES code can be downloaded at http://www.astro.up.pt/ determine the EW by integrating the line

5 Table 3 Stellar parameters and metallicities from the UVES and UES spectra

HD Teﬀ δTeﬀ log g δ log g ξt δξt [Fe/H] δ[Fe/H] [K] [K] [dex] [dex] [km s−1] [km s−1] [dex] [dex]

106252 5880 15 4.40 0.02 1.13 0.02 -0.070 0.011 117207 5680 29 4.34 0.04 1.06 0.04 0.250 0.022 12661a 5760 37 4.33 0.07 1.09 0.04 0.385 0.028 216437 5882 21 4.25 0.03 1.25 0.02 0.250 0.018 217107a 5679 42 4.32 0.07 1.15 0.05 0.350 0.034 28185 5726 40 4.45 0.05 1.08 0.05 0.230 0.031 4203 5728 49 4.23 0.07 1.18 0.06 0.430 0.036 70642 5732 24 4.42 0.06 1.06 0.03 0.190 0.018 73526 5666 25 4.17 0.04 1.12 0.03 0.260 0.020 76700 5694 27 4.18 0.05 1.05 0.03 0.370 0.023

2 Note.—Stellar parameters, Teﬀ and log(g/cm s ), microturbulent velocities, ξt, and metallicities , [Fe/H], and their uncertainties, of the planet-host stars observed with UVES/VLT and UES/WHT spectrographs. aStars observed with the UES spectrograph at WHT telescope.

flux; whereas for Zr, we “manually” per- a factor of 1.5 times the rms. We also formed a gaussian fit taking into account checked that we get the same results when possible blends. In both cases, we use the using a factor of 2 times the rms, but task splot within the IRAF package. Fi- we decided to stay in a restrictive posi- nally, the EWs of Sr, Ba and Zn lines were tion. The line-by-line scatter in the dif- checked within IRAF, because, as well as ferential abundances goes on average, for for the elements O, S, Eu and Zr, for some the 5 “single” solar twins (see Sect. 6.2), stars, ARES did not find a good fit due to from σ ∼ 0.012 − 0.014 for elements like numerical problems and/or bad continuum Ni, Si and Cr, to σ ∼ 0.048 − 0.056 for location. For further details see Sect. 5.2. Mg, Zn and Na, whereas for the whole Once the EWs are measured, we use sample of 71 solar analogs without plan- the LTE code MOOG (Sneden 1973) to ets, the line-by-line scatter goes on average compute the chemical abundance provided from σ ∼ 0.016 for elements like Ni, Si and by each spectral line, using the appro- Cr, to σ ∼ 0.050 for Mg and Na and to priate ATLAS model atmosphere of each σ ∼ 0.074 for Zn. star. We determine the mean abundance of each element relative to its solar abun- 5.1. Atomic data dance (see Sect. 5.2) by computing the The oscilator strengths of spectral lines line-by-line mean difference. However, to used in this study were extracted from avoid problems with wrong EW measure- three different works. The line data from ments of some spectral lines, we rejected Na (Z = 11) to Ni (Z = 28) were compiled all the lines with an abundance different from Neves et al. (2009). Among these el- from the mean abundance by more than ements, the only element with abundance

6 derived from an ionized state is Sc. C I and be mentioned that the stars We note here S I data is based on Ecuvillon et al. (2004), that the spectral lines in solar spectra ob- and [O I] from Ecuvillon et al. (2006a). tained on the daytime sky are known to We note that S I multiplet lines were com- exhibit EW and line depth changes (e.g. bined into a single line by adding the gf Gray et al. 2000). Therefore, although values. The three Zn I were extracted from the asteroid spectrum has a lower S/N ra- Ecuvillon et al. (2004) and Reddy et al. tio than the sky spectrum, we think it is (2003), as well as the four Cu I lines. How- more likely a better reference solar spec- ever, three Cu I lines were rejected high trum and we have adopted the asteroid dispersion in the resulting Cu abundances element abundances as our solar reference probably due to blending effects with other in this work. However, we will compare element lines. The atomic data of the s- the results obtained with the sky spec- process elements Sr, Ba, Y, Zr, Ce, Nd trum when analysing the solar twins (see and the r-process element Eu were also ex- Sect. 6.2), just to show how important is tracted from Reddy et al. (2003). Simi- to have a good solar reference spectrum. larly, one Ba I line was discarded since its It should be mentioned, however, that the too high strength provided very different UVES and UES stars were also analysed abundance from the other two lines (see using the HARPS Ganymede spectrum as Sect. 5.2). solar reference due to the lack of an appro- We perform a completely differential priate solar spectrum observed with these analysis to the solar abundances on a line- two instruments. by-line basis, and therefore, the uncertain- 0.15 ties on the oscilator strengths are nearly O irrelevant, since most of the lines are in 0.10 O Zr Eu the linear part of the curve of growth, with Sr Y S Nd C Mg Zn Ba Mn Eu C Na Sr the exception of some Ba, Fe, and probably 0.05 Sc Cr Cu Na Fe Cu Ce V Ni Zr MgAlSi S CaTi Co ATLAS-HARPS Cr Co Zn BaCe also some Zn, Mn, and Ca lines. Sc V Si Ti Mn Y Ni Ca Fe Nd

[X/H] 0.00 5.2. The solar reference ∆ -0.05 This fully differential analysis is, at Sky solar spectrum S/N = 1000 Ganymede solar spectrum S/N = 400 least, internally consistent. For this rea- -0.10 0 10 20 30 40 50 60 70 son, we have used two HARPS solar Z spectra3 as solar reference: a daytime sky spectrum (Dall et al. 2006) with a Fig. 2.— Abundance difference on a line- S/N ∼ 1000 − 1200 and the spectrum of by-line basis between the Kurucz atlas so- the Ganymede, a Jupiter’s satellite, with lar spectrum and, a daytime sky spectrum a S/N ∼ 350 − 400. However, it should (filled circles) and a Ganymede spectrum (open circles). The error bars show the 3The HARPS solar spectra can be downloaded at standard deviation of the line abundance http://www.eso.org/sci/facilities/lasilla/instruments differences divided by the square root of /harps/inst/monitoring/sun.html the number of lines of each element.

7 In Fig. 2 we display the difference be- al. 2002). In the atlas solar spec- tween the solar chemical abundances mea- trum the whole feature has an EW sured in the Kurucz atlas solar spectrum of ∼ 5.4mA˚ which provides an abun- (Kurucz et al. 1984) and the sky and dance, A(O)4, of 8.74 dex. It is the Ganymede solar HARPS spectra, com- element with the smaller EW mea- puted in the same way as the chemical surement in this study. abundances of the whole sample of solar analogs. The abundances of the elements b) Zirconium abundances come from ˚ O, Sr, Cu, Eu, Nd, and Zr were determined the line Zr II λ 5112.28 A which from only one spectral line; those of the el- slightly blended with two weaker ∼ ements C, S, Na, Mg, Al and Ba from 2 lines. We estimate an EW of ˚ lines; those of Zn, Y, Ce from three lines; 9.5 mA in the solar atlas spectrum. and the abundances of the rest of elements In the HARPS spectrum of the as- in Fig. 2 from more than 5 lines. In Fig. 2 teroid, the Zr II abundance is sub- we adopt an uncertainty of 0.04 dex and stantially different from that of the 0.015 dex for the asteroid and sky spec- atlas spectrum. In addition, the sky trum, respectively, for the elements with and the asteroid spectra show dif- ∼ ˚ ˚ only one line. For the other elements, the ferent EW, 9.2 mA and 8.6 mA, adopted uncertainty is computed from the respectively, which translates into an ∼ standard deviations of the individual line abundance difference of 0.05 dex. measurements divided by the square root This does not seems to be related of the number of lines. Most of the ele- with the S/N ratio, which provides ment abundance differences with respect an uncertainty on the EW measure- ∼ ˚ ˚ to the Sky HARPS spectrum are slightly ment of 0.05 mA and 0.15 mA, higher than zero, whereas these abundance respectively, from the Cayrel’s for- differences are roughly zero in the case of mula (Cayrel 1988). the Ganymede HARPS spectrum. This is c) Neodymium is estimated from the probably related to the EW variations in line Nd II λ 5092.80 A˚ which is rel- the sky daytime solar HARPS spectrum. atively isolated or weakly blended. In general, elements with relatively small However, in this case the sky spec- ˚ line equivalent widths (EW . 10 mA) trum clearly show smaller EW than show the larger discrepancies with the so- the asteroid spectra, ∼ 7.6 mA˚ and lar atlas abundances. We list below some 9.2 mA,˚ respectively, resulting into problem- detailed information on the most an abundance difference of ∼ 0.1 dex. atic elements: d) Europium is determined from the a) Oxygen abundance is derived from line Eu II λ 6645.13 A,˚ and is also the forbidden [O I] line at 6300.304 A˚ blended with two lines of smaller (with a expected EW of ∼ 4.1 mA)˚ EWs. The solar atlas spectrum and is severely blended with the shows EW of ∼ 6.4 mA.˚ [Ni I] λ 6300.34 A˚ (with a expected 4 EW of ∼ 1.3 mA,˚ e.g. Nissen et A(X)= log[N(X)/N(H)] + 12

8 Sr + + Planet-host: 2 stars |S/N| = 810 − 166 Sr Planet-host: 2 stars |S/N| = 810 − 166 0.10 0.10 S C S Single: 5 stars |S/N| = 907 −+ 448 Single: 5 stars Co |S/N| = 722 −+ 456 Nd Zr Y Mn Mn C Na Sr Zr S Mn S O Co Mn CoNi Y Sr Na Ni Nd 0.05 Na 0.05 Na Cu Ni V Cu VCe MgNi C Zn Cr C Cu V Si Y O Zn Al Nd

SUN-STARS V SUN-STARS SiCo Cu Cr Si Y Zr SiEu Ca CrFe Ti Sc FeEu Ti 0.00 Eu Ba Al 0.00 Ti Mg Ba O BaCa NdAl Eu Ce Ti Mg Ba [X/Fe] Sc [X/Fe] Ca Sc ∆ Zr ∆ Mg

-0.05 -0.05 O

TEFF-TSUN = 77 K Logg-LoggSUN = 0.10 TEFF-TSUN = 77 K Logg-LoggSUN = 0.10 S/N > 370 -0.07 < [Fe/H] < 0.07 S/N > 370 -0.07 < [Fe/H] < 0.07 -0.10 -0.10 0 500 1000 1500 0 500 1000 1500

TC (K) TC (K)

Fig. 3.— Mean abundance differences, ∆[X/Fe]SUN−STARS, between the Sun, and 2 planet hosts (filled circles) and 5 “single” stars (open circles). Error bars are the standard deviation from the mean divided by the square root of the number of stars. Linear fits to the data points weighted with the error bars are also displayed for planet hosts (solid line) and “single” stars (dashed line). The left panel shows the results using the solar spectrum of the Ganymede and the right panel, using the sky solar spectrum taken in daytime. Note that the average S/N ratios are different in both panel for “single” stars. This is due to the fact that different stars fulfill the [Fe/H] condition of solar twins (see Sect. 6.2).

e) Barium, which comes from the two three lines Y II at 5087.43, 5200.42, strong lines of Ba II λ 5853.69, and 5402.78 A.˚ As Zr, Ba and Ce, and λ 6496.91 A,˚ can be easily mea- these two elements also show a large sured, in principle, although the dispersion in their Galactic chemical [Ba II/Fe] versus [Fe/H] trend dis- evolution trends (see Sect. 6.1). plays the largest dispersion among the elements analyzed in this work 6. Discussion (see Sect. 6.1). This may be related to the dependence of the strong In this section we will discuss the abun- Ba II lines on the microturbulence dance ratios of different elements as a func- and other factors. tion of the metallicity, [Fe/H], as well as the relation, for different samples of so- f) Cerium is derived from three lines lar analogs, between the abundance differ- of Ce II at 4523.08, 4628.16, and ence, ∆[X/Fe]SUN−STARS and the 50% equi- ˚ 4773.96 A. The error bars are rela- librium condensation temperature, TC , for tively small but it is also off from the a solar-system composition gas with 50% zero level in Fig. 2. element condensation, which is the temperature at which half of an element is kept g) Strontium and Yttrium abundances in the gas and the other half is kept into are estimated from the relatively condensates (Lodders 2003). strong lines Sr I λ 4607.34 A˚ and the

9 6.1. Galactic abundance trends by their error bars. These fits show the decreasing trend of the mean abundance The high-quality of HARPS and UVES ratios with condensation temperature al- data presented in this work allows us to ready reported in Meléndez et al. (2009) derive very accurate abundance ratios of but the trend is not so clear. We note many elements. In Figs. 11, 12 and 13 here that the S/N is greater than 500 for (available online), we show the abundance all stars in this figure, except one “single” ratios [X/Fe] versus [Fe/H] for the whole star with a S/N ∼ 370. We display on the sample of solar analogs analyzed in this upper-right corner of both panels of Fig. 3 work. These trends are compatible with the mean S/N and its standard deviation. those in previous works (e.g. Bodaghee et Besides, the mean S/N ratios are different al. 2003; Beirão et al. 2005; Bensby et al. in both panel for “single” stars. This is 2005; Gilli et al. 2006; Reddy et al. 2006; due to the fact that just one star in each Takeda 2007; Neves et al. 2009; Ram´ırez panel did not fulfill the [Fe/H] condition of et al. 2009). We note here the low disper- solar twins. This is because small offsets sion of most of the elements analyzed in in the range 0.01-0.05 dex are found in the this work, with the exception of some ele- differential Fe abundances and probably ments discussed in Sect. 5.2. In Table 4 we different Fe lines were used for each star in provide an example of the tables contain- each case (see Sect. 5). In the right panel ing all the element abundance ratios [X/Fe] of Fig. 3 the fits change completely mainly which are all available online. due to the position of O and Zr abundance 6.2. HARPS solar twins ratios. We have checked that expanding the metallicity range up to ±0.15 dex only We examined the whole sample of increases the number “single” stars, up to HARPS targets trying to search for so- 11, which is the same number of stars as in lar twins defined in the same way as in Meléndez et al. (2009). However, this does Meléndez et al. (2009) and we found 2 not change significantly either the position planet host and 5 “single” stars (see Ta- of the data points or the slope of the linear ble 2). In Fig. 3 we display the mean abun- fit. dance difference, ∆[X/Fe]SUN−STARS, ver- We find some slight differences between sus the condensation temperature, TC , for the abundance ratios of stars hosting plan- both reference solar spectra, the Ganymede ets and “single” stars. However, the fits spectrum (left panel) and the daytime sky have similar slopes, but only in the left spectrum (right panel). The main differ- panel of Fig. 3. In addition, the number ences between both panels are for the el- of planet hosts is too small to allow us to ements O, Nd, Zr, and Y which were dis- make a strong statement on implications cussed in Sect. 5.2. Although the error bars of these differences. are relatively large to search for any clear We have tried to understand why our trend, it seems that the mean abundance results do not look like those of Meléndez ratios of refratories are on average smaller et al. (2009). Our data has better S/N and than those of volatiles. In this plot we also resolving power than those of Meléndez et show linear fits to the data points weighted

10 Table 4: Abundance ratios [X/Fe] of solar analogs without known planetsa

HD [C/Fe] [O/Fe] [S/Fe] [Na/Fe] [Mg/Fe] [Al/Fe]

Mel´endez et al. (2009), according to our 0.10 C S Single: 5 stars error bars.

O Mn 0.05 Na

Zn MgNi Cu Y 1.0 Co V Al

SUN-STARS Si Cr 0.00 Ti Sc Ba 0 < TC (K) < 1800

[X/Fe] Ca 0.8 TC (K) < 1200 ∆ Zr

TC (K) > 1200 -0.05 0.6

TEFF-TSUN = 77 K Logg-LoggSUN = 0.10 S/N > 350 ∆[Fe/H] = 0.07 -0.10 Confidence 0.4 0 500 1000 1500 T (K) C 0.2

Fig. 4.— 100 Monte Carlo simulations of 0.0 -1.5 -1.0 -0.5 0.0 0.5 T slope/10-4 the mean element abundance ratios (aster- C isks) reported in Meléndez et al.(2009) taking into account the error bars shown in Fig. 5.— Histograms of slopes of linear fits Fig. 3, and in comparison with the results to 1000 simulations of the points on the (filled circles) shown in Fig. 3. fits to mean element abundance ratios reported in Meléndez et al. (2009). The histograms show the distributions for TC [K] < al. but it seems that our dispersion in 1200 (dashed line), TC [K] > 1200 (dashed- the mean abundance ratios versus TC and dotted line), and 0 < TC [K] < 1800 (solid also around the linear fits is larger than in line). On the top of this figure, we display Meléndez et al. (2009). The number of so- the slopes and error bars of our results for lar twins in our study is smaller, although the 5 “single” twins stars shown in Fig. 3. we think this may not explain the differ- See also Table 5. ences between both studies. We have made two Monte Carlo simulations to check if In the first simulation (see Fig. 4), our results are consistent with those of we randomly generate 5 abundance ratios

11 (corresponding to the 5 solar twins) using TEFF-TSUN= 177 K Logg-LoggSUN=0.3 gaussian distributions around the mean 0.1 S/N > 370 -0.30 < [Fe/H] < 0.50 S Zn Nd Sr MgCoNi Ca Mn SiFe abundance ratios reported in Meléndez et O Y C Cr Cu BaCe Ti Sc Na Al -0.0 V Zr al. (2009) taking into account the error Eu Single: 71 stars + Nd bars displayed in the left panel of Fig. 3. SUN-STARS O |S/N| = 841 − 449 Si Cr -0.1 C Sr Ca We then compute the mean abundance Fe Ce Ti [X/H] S Mn MgCoNi ∆ Ba Y Zn Eu Sc Zr V ratios from these 5 points and finally dis- Na Cu Al played 100 simulations for each element -0.2 Planet-host: 24 stars in Fig. 4 in comparison with our results |S/N| =1039 −+ 359 -0.3 of the 5 “single” stars (see left panel of 0 500 1000 1500 T (K) Fig. 3). Note that Meléndez et al. (2009) C do not include in their analysis the elements Ce, Nd, Sr and Eu, which are in Fig. 6.— Same as in Fig. 3 but for the fact some of the outliers in Fig. 2 (see also mean ratio ∆[X/H]SUN−STARS in the 24 Sect. 5.2), although they include N, P, and planet host and 71 “single” stars of the K. This simulation shows that our results whole sample of solar analogs, using the and those in Meléndez et al. are indeed Ganymede spectrum as solar reference. consistent, although in some cases only marginally, for most of the elements, ac- In the second simulation, we derive a cording to the error bars shown in Fig. 3, linear fit to the mean abundance ratios with the exception of Si, Ca and Cu. The presented in Meléndez et al. (2009), and line-by line scatter of Ca and Si differential define, as new reference abundance ratios, abundances for the 5 “single” stars are on the points on this fit at the condensation average 0.013 and 0.018 dex whereas the temperature of each element. Then we star-to-star scatter, 0.022 and 0.019 dex, generate, around these new points, abun- respectively. The star-to-star scatter for dance ratios using again gaussian distribu- Cu is 0.027 dex. Note that the error bars tions which take into account the error bars in Fig. 3 contain the star-to-star scatter di- of the results of “single” stars (see Fig. 3). vided by the square root of the number of Finally, for each simulation we perform a stars, but only one line of Cu was used in linear fit to these abundance ratios. We re- the abundance computation (see Sect. 5.1) peat the same procedure only for elements The fact that the star-to-star and line-by- with TC above or below 1200 K. Three his- line scatters are roughly equal and so small tograms of 1000 simulations are displayed may exacerbate the disagreement between in Fig. 5. The slopes of the linear fits com- our Ca and Si mean abundances and those ing from our data are shown on the top of in Meléndez et al. (2009). In addition, the this figure (see also Table 5) and seems to scatter of our mean element abundance ra- be clearly consistent with the peak of these tios around the linear fits is higher than in histograms, although for TC [K] > 1200, Meléndez et al., which may be related to our point slightly moves to higher values the solar reference HARPS spectrum used of the slopes. In Table 5 we provide the in this analysis. values of the slopes in the linear fits to the

12 TEFF-TSUN = 177 K Single: 71 stars S/N > 370 All planet-host: 24 stars 0.2 Planet-host: 10 stars with 1000 < PORB < 4300 days Logg-LoggSUN = 0.3 Planet-host: 10 stars with 150 < P < 650 days -0.30 < [Fe/H] < 0.50 ORB Planet-host: 4 stars with PORB < 25 days S Nd C Sr Mn O Na Ni Y Zn Cu CrSiMgCo V Ca Al Zr 0.0 Eu BaCe Ti Sc O Nd C S Ba Sr Y Zn CrMgEu Ca Zr SiNi Ti SUN-STARS O Na Cu Mn Co VCe NdAl -0.2 Ba Sc C S Sr Eu Y Zn CrMg Ca Zr Mn SiNi Ce Ti [X/Fe] V Al O Na Cu Co Sc ∆ C S Ba Nd Sr Y -0.4 Zn MgEu Ca O Na CrSiNi Cu Mn Ce Ti Al Zr Co V NdSc C S Zn Ba Sr Zr MgEu Y Na CrSi CeCa Ti -0.6 Cu Ni V Al Mn Co Sc

0 500 1000 1500

TC (K)

Fig. 7.— Same as in Fig. 3 containing 71 “single” stars (open circles) and 24 stars hosting planets (filled circles), with the most massive planet in an orbital period Porb < 25 days (4 stars, filled diamonds), 150 < Porb < 650 days (10 stars, filled squares), 1000 < Porb < 4300 days (10 stars, filled triangles). The ∆[X/H]SUN−STARS and linear fits has been artificially shifted by −0.15 dex, for the sake of clarity. Horizontal dashed-dot lines show the zero point levels for each set of points.

mean abundance ratios ∆[X/Fe]SUN−STARS (see Fig. 7), although due to higher mean as a function of the condensation temper- metallicity of the stars hosting planets, atures TC , at different ranges of TC . their mean abundance ratios appear shifted. In addition, the error bars in this plot are 6.3. All solar analogs larger than in a plot ∆[X/Fe]SUN−STARS, since the dispersion in [X/H] is typically The whole sample of solar analogs pre- larger than in [X/Fe] due to chemical evo- sented in this paper contains 24 planet lution effects. Those elements, like Mn hosts and 71 stars without known plan- and O, with steeper trends [X/Fe] versus ets. In Fig. 6 we display ∆[X/H] SUN−STARS [Fe/H], will show larger abundance differ- of the sample versus T , only using the C ences in the plot ∆[X/Fe] (see Ganymede spectrum as solar reference. SUN−STARS Fig. 7). However, in the plot ∆[X/H] The abundance pattern in this figure is SUN−STARS these differences are relatively smaller and very similar to that of ∆[X/Fe] SUN−STARS most points seems to agree well with the

13 Table 5

Slopes of the linear fit to the mean values ∆[X/Fe]SUN−STARS versus TC

a b Sample 0 1200 Nstars

sSA −0.358 ± 0.031 −0.446 ± 0.066 −0.487 ± 0.070 71 pSA −0.322 ± 0.049 −1.161 ± 0.113 −0.286 ± 0.103 24 smrSA −0.163 ± 0.052 −1.158 ± 0.113 0.188 ± 0.115 14 pmrSA −0.138 ± 0.046 −1.451 ± 0.100 0.327 ± 0.104 14 sST −0.449 ± 0.089 −0.464 ± 0.210 −0.307 ± 0.202 5 pST −0.585 ± 0.025 0.198 ± 0.147 −0.644 ± 0.028 2 PHs 0.075 ± 0.075 −1.346 ± 0.183 0.298 ± 0.123 4 PHm −0.548 ± 0.080 −1.093 ± 0.196 −0.623 ± 0.167 10 PHl 0.086 ± 0.066 −1.109 ± 0.154 0.154 ± 0.144 10

Note.—Slopes of the linear fit of the mean abundance ratios, ∆[X/Fe]SUN−STARS, as a function of the condensation temperature, TC , using the elements with TC within the specified interval, for different stellar samples. aStellar samples: “single” solar analogs, “sSA”, planet-host solar analogs, “pSA”, “single” metal-rich solar analogs, “smrSA”, planet-host metal-rich solar analogs, “pmrSA”, “single” solar twins “sST”, planet-host solar twins “pST”, planet-host solar analogs with the most massive planet in an orbital period Porb < 25 days, “PHs”, 150 < Porb < 650 days , “PHm”, 1000 < Porb < 4300 days, “PHl”. bTotal number of stars including those with and without known planets. linear fits. Both stars with and with-

out planets show a similar decreasing 10.00 10.00 trends towards increasing values of TC . ]

Nearly equal results are found for the plot Jup

1.00 1.00 ∆[X/Fe]SUN−STARS versus TC . Planet Mass [M Jup ] The stars hosting planets studied in this Planet Mass [M work shows a variety of planetary systems 0.10 0.10 with diﬀerent orbital periods, Porb, from 3 to 4200 days, and minimum masses, i.e.

0.01 0.01 Mp sin i, from 0.02 to 17.4 MJup, where Mp 10 100 1000 -0.2 0.0 0.2 0.4 is the mass of the primary planet, i.e. the Orbital Period [days] [Fe/H] most massive planet in the planetary sys- Fig. 8.— Mass of the most massive planet tem, M , the Jupiter mass, and i, the Jup of each star versus the orbital period of the orbital inclination. Among the 24 planet- planet (left panel) and metallicity [Fe/H] host stars in the sample, six are plane- (right panel). The symbols as in Fig. 7. tary systems containing two planets, and one of these planetary systems, in the star HD 160691 (µ Arae, e.g. Pepe et al. 2007), ing the smallest a super-Earth like planet contains four planets known so far, be- with a minimum mass of ∼ 10.5 Earth

14 masses and its primary planet a Jupiter- Although the linear fits in Fig. 7 to the like planet with Mp sin i ∼ 1.8 MJup (see element abundance ratios seem to change the extrasolar-planet encyclopaedia5). We with the orbital period, the position of do not really know the nature of this and most of these abundance ratios remains the other super-Earth like planets and they same and only few elements like Zn, Nd may very well be Neptune like ice giants and Zr modify significantly their position (see e.g. Barnes et al. 2009). with respect to other elements and thus There is only one star, HD 202206, con- changing the slope of these fits. We sus- taining a planet with Mp sin i > 10 MJup, pect that these changes may be linked to which has a primary planet with a Mp sin i = chemical evolution and abundance scatter, 17.4 MJup at Porb ∼ 256 d, and a sec- and hardly related to the different orbital ondary planet with a Mp sin i = 2.44 MJup periods of the primary planets. In Table 5 at Porb ∼ 1383 d (e.g. Udry et al. 2002; we give the slopes for the sub-samples of Correia et al. 2005). This primary planet planet-host stars at different ranges of or- may be considered as a brown dwarf, bital periods. One can appreciate that, due whose minimum mass may be settled at to this large scatter in the mean abundance ∼ 13 MJup, and/or entering in the so-called ratios, the derived slopes are very different brown dwarf desert (e.g. Halbwachs et al. for these different sub-samples. However, 2000). the number of stars in these different sub- In Fig. 7 we display the mean abun- samples is so small that we cannot extract any clear conclusion from this separation dance ratios ∆[X/Fe]SUN−STARS for the 24 stars hosting planets according to the or- in three orbital period intervals. bital period of the primary planet in the 6.4. Metal-rich solar analogs planetary system of each star. We have separated the sample in three ranges: a) 4 To evaluate the abundance trends with stars with primary planets at Porb[d] < 25; the same number of planet hosts and “sin- b) 10 stars at 150 < Porb[d] < 650; and gle” stars, we studied a super-solar metal- c) 10 stars at 1000 < Porb[d] < 4300. We licity sample of solar analogs with Teff − note that Jupiter has an orbital period of Teff,⊙ = ±177 K and log g − log g⊙ = 4333 days around the center-of-mass of the ±0.2 dex and 0.15 . [Fe/H] . 0.4. This solar system. The distribution of masses of sample contains 14 stars with and 14 with- the most massive planet in each planetary out planets. In Fig. 9, we show the trends system as function of orbital periods and ∆[X/Fe]SUN−STARS versus TC . The posi- metallicities are shown in Fig. 8. There is tion of some elements in this plot is cer- no case in this sample in which the most tainly affected by chemical evolution ef- massive planet has orbital period in the fects due to the high metal content of the gap in between these ranges. Of course, sample, in particular , Mn and O. This may secondary planets in these systems may lie explain the higher scatter of the points in at orbital period within these gaps. this plot with respect to the linear fits. However, what is relevant from this plot is 5http://exoplanet.eu that both samples of stars with and with-

15 out planets show almost exactly the same ranges, 4700 . Teff [K] . 6400 and −0.8 . abundance pattern. All the element abun- [Fe/H] . 0.5. They did not find any dance ratios consistent in both samples clear difference between these two sam- within the error bars, except for the Zn, ples. They also searched for dependence Mn, Ba, and Nd, which still are very close on the planet orbital parameters: mass, or- in this plot. In addition, the linear fits have bital period, eccentricity and orbital sepa- almost exactly the same slope which agrees ration without any success. Ram´ırez et al. with the previous statement (see also Ta- (2009) have derived the slopes in [X/Fe] ble 5). It should be mention that among for two TC ranges above and below 900 K the stars hosting planets, there is a vari- in a sample of solar analogs and twins. ety of planetary systems (see Fig. 8) with There appears to be two distinct groups planets of different masses at different or- for super-solar metallicities, showing pos- bital periods. In Sect. 6.5 we provide more itive and negative slopes. Following their details on these planetary systems and dis- line of reasoning, a null (solar-like) or even cuss the slopes of these trends for each negative slope implies that a great fraction planet-host and “single” star, in the con- of refractory elements have been extracted text of the presence of terrestial planets. from the star-forming cloud to make up dust grains, also suggesting planet forma-

0.15 O tion. Thus, they tentatively conclude that O TEFF-TSUN = 177 K Logg-LoggSUN = 0.2 S/N > 500 0.14 < [Fe/H] < 0.36 Nd this bimodal distribution of slopes is sep- 0.10 Ba Nd arating super-solar metallicity stars with C Zn Ba S S and without terrestial planets. Unfortu- 0.05 Ca Eu Ca SUN-STARS Zr Mg Y nately, they do not know how many of their Zn Zr Cr CrFe 0.00 Sr SiMgEu Ti stars have planets, whatever their masses [X/Fe] Ce

∆ Al Ni Cu Al Ni VCe Planet-host: 14 stars Cu Co Sc are. One should also note that the core of a V -0.05 |S/N| =1138 −+ 410 Mn Na Co Sc Jupiter-like planet should contain the same Na Single: 14 stars Mn |S/N| = 848 −+ 254 amount refractory elements than roughly -0.10 0 500 1000 1500 three or four terrestial planets. However, T (K) C Jupiter-like planets also have a substantial amount of volatiles, whereas terrestial Fig. 9.— Same as in Fig. 3 for 14 planet planets do trap only refractories. host and 14 “single” stars of the sam- We have determined the slopes of all so- ple of metal-rich solar analogs, using the lar analogs of our sample with and with- Ganymede spectrum as solar reference. out known planets. In Fig. 10, we display the slopes of the trends ∆[X/Fe]SUN−STARS 6.5. Planetary signatures in the versus the whole range of TC for all the abundance trends? stars in our sample. We note that the slopes in our study have the same sign Ecuvillon et al. (2006b) measured the than in Mel´endez et al. (2009) but oppo- slopes of the trends [X/H] versus TC in site than in Ram´ırez et al. (2009), because a sample of 88 planet-host stars and 33 the later used the abundance ratios [X/Fe] “single” stars in a large Teﬀ and [Fe/H]

16 et al., this result also implies that most 14 of the stars with and without planets in 12 our sample would not have terrestial plan- 10 Single stars ets. However, in two of them, it has been Planet hosts 8 already detected super-Earth like planets ∼ − 6 with masses in the range 7 11 Earth Number of stars masses. The most massive one is in the 4 planetary system of the star HD 160691, 2 with a Jupiter-like planet at Porb ∼ 4205 d 0 -1.5 -1.0 -0.5 0.0 0.5 and the other one is an isolated planet T slope/10-4 C around the star HD 1461 (e.g. Rivera et al. 2010, see Fig. 8). These two planet- Fig. 10.— Histograms of slopes of lin- host stars show clearly a negative slope ear ﬁts to 71 “single” stars (solid line) with –0.58 and –0.26 ×10−4 dex/K which and 24 stars hosting planets (dashed line) is exactly the opposite to what is expected of the mean element abundance ratios, from Ram´ırez et al. (2009). Positive ∆[X/Fe]SUN−STARS, versus the condensa- slopes or consistent with zero within the tion temperature, TC . The mean error bars error bars are found for planet hosts which of the slopes are shown in upper-right cor- have either only one Jupiter-like planet at ner. Porb & 1200 d or three planetary systems with two planets. Two of them, HD 217107 instead of ∆[X/Fe]SUN−STARS. Although and HD 12661 with two Jupiter-like plan- the number of planet hosts is smaller than ets each, at very diﬀerent orbital periods “single” stars, the peak of the distribu- (e.g. Wright et al. 2009). The other tion of the slopes of these two samples is one, HD 47186, contains a Neptune- and a centered around the same position, which Saturn-like planets at Porb ∼ 4 and 1354 d, corresponds, in fact, to a negative slope. respectively (Bouchy et al. 2009). We may We may refer here to the Table 5 where conclude that there is no reason to expect we provide the slope of the mean abun- that these stars hosting relatively massive dance ratios, ∆[X/Fe]SUN−STARS, in the planets should also contain terrestial plan- whole range of TC of the all solar analogs of ets while the other stars with a already de- our sample with and without known plan- tected super-Earth like planet should not, ets. These values are slightly higher than and/or that the amount of refractory met- the position of the peaks in Fig. 10 but als in the planet hosts depends only on the the values of the slopes are very similar amount of terrestial planets. In addition, and consistent with the error bars for both it seems plausible that many of our targets stars with and without planets. This in- hosts terrestrial planets. This is supported dicates that stars with already detected by numerical simulations which show that planets behave in similar way, with respect terrestrial planets are much more common to the chemical abundances, as stars with- than giant planets, and probably 80–90% out known planets. In addition, accord- of solar-type stars have terrestrial planets ing to the line of reasoning in Ram´ırez (e.g. Mordasini et al. 2009a,b). This state-

17 ment agrees with the growing population lar analogs in the range [Fe/H] = 0.25 ± of low mass planets found in the HARPS 0.11 dex. This sample has 14 stars host- sample of exoplanets (see e.g. Udry & San- ing planets and 14 “single” stars. In this tos 2007). case, the abundance pattern and the slope of the linear fit to both samples are al- 7. Conclusions most equal, which may indicate that the abundance pattern found by Meléndez et With the aim of investigating possible al. (2009) do not have nothing to do with connection between the abundance pat- the presence of planets. tern versus the condensation temperature, Finally, we derive the slopes of the trend T , and the presence of terrestial planets, C ∆[X/Fe] versus T for each star we have selected a sample of solar analogs SUN−STARS C of the whole sample and compare the dis- with and without planets from very high tribution of the slopes found for stars with signal-to-noise and very high resolution and without planets. Although the num- HARPS, UVES and UES spectra. The ber of planet hosts is significantly smaller whole sample contains 71 “single” stars than that of “single” stars, the peaks of the and 24 planet hosts with 5600 < T [K] < eff distributions are placed at the same posi- 5954, 4.0 < log(g[cm s−2]) < 4.6, −0.3 < tion, at about −0.5 × 10−4 dex/K, which [Fe/H] < 0.5. We perform a detailed chem- is exactly the opposite sign than that re- ical abundance analysis of this sample and quired by the presence of terrestial planets, investigate possible trends of mean abun- following the line of reasoning in Ram´ırez dance ratios, ∆[X/Fe] , versus SUN−STARS et al. (2009). Thus, according to Ram´ırez T . We find that both stars with and C et al. (2009), most of our planet-host without planets show a similar abundance stars would not contain terrestial planets, pattern, except for some elements like Mn a statement that a priori does not appear which are affected by chemical evolution to be expected. effects, due to the higher metal content of the stars hosting planets. J.I.G.H. acknowledges financial sup- Only in the sample of HARPS spectra, port from the Spanish Ministry project there are 7 solar twins, 2 of them hosting MICINN AYA2008-00695 and also from planets. We study the mean abundance the Spanish Ministry of Science and In- ratios as a function of T and compare C novation (MICINN) under the 2009 Juan them with the results of Meléndez et al. de la Cierva Programme. N.C.S., S.S. and (2009). Our results are consistent with V.N. would like to thank the support by their results within our error bars, which the European Research Council/European has been demonstrated through two Monte Community under the FP7 through a Carlo simulations. Nevertheless, the scat- Starting Grant, as well from Funda¸cão ter around the linear fits in our data is para a Ciência e a Tecnologia (FCT), larger, which may be related to the solar Portugal, through a Ciência 2007 contract reference HARPS spectrum used in this funded by FCT/MCTES (Portugal) and analysis. POPH/FSE (EC), and in the form of grant We also select a metal-rich sample of so-

18 reference PTDC/CTE-AST/098528/2008 Ecuvillon, A., Israelian, G., Santos, N. C., from FCT/MCTES. E.D.M, J.I.G.H. and Mayor, M., Villar, V., & Bihain, G. G.I. would like to thank ﬁnancial sup- 2004, A&A, 426, 619 port from the Spanish Ministry project MICINN AYA2008-04874. S.S. acknowl- Ecuvillon, A., Israelian, G., & Santos, N. edges the support of the FCT grant: C. 2006, A&A, 445, 633 SFRH/BPD/47611/2008. This work has Ecuvillon, A., Israelian, G., Santos, N. C., also made use of the IRAF facility, and Mayor, M., & Gilli, G. 2006, A&A, 449, the Encyclopaedia of extrasolar planets. 809

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21 0.2 0.2

0.0 0.0 [C/Fe] [O/Fe] -0.2 -0.2