UNIVERSITE´ DE GENEVE` FACULTE´ DES SCIENCES D´epartement d’Astronomie Professeur Michel Mayor Professeur St´ephane Udry

Metallicity-biased search for transiting Hot

THESE`

pr´esent´ee`ala Facult´edes Sciences de l’Universit´ede Gen`eve pour obtenir le grade de Docteur `esSciences, mention Astronomie et Astrophysique

par Ronaldo Oliveira DA SILVA

de

S˜ao Paulo (Br´esil)

Th`eseNo 3941

GENEVE` Observatoire de Gen`eve 2008

Ce travail de th`ese a pu ˆetre realis´egrˆace au support financier accord´epar Coordena¸c˜ao de Aperfei¸coamento de Pessoal de N´ıvel Superior (CAPES, Br´esil), sous la forme d’une bourse d’´etudes. Tous mes sinc`eres remerciements.

R´esum´een fran¸cais

Apr`es la d´ecouverte de 51 Peg b par Mayor & Queloz (1995), la premi`ere plan`ete au- tour d’une ´etoile autre que le Soleil, plusieurs auteurs ont plac´eleurs efforts `ala recherche de nouveaux mondes. Les caract´eristiques inattendues de 51 Peg b, une grosse plan`ete (0.47 MJup) en orbite `atr`es courte p´eriode (4.2 jours seulement), ont donn´eencore plus de motivation aux recherches, toujours en se posant la question de pourquoi ces plan`etes sont aussi diff´erentes de celles de notre syst`eme solaire. Jusqu’`apr´esent, environ 250 de ces plan`etes extrasolaires ont ´et´ed´ecouvertes, avec des allant de quelques masses ter- restres jusqu’`aplusieurs masses de , distances orbitales allant de quelques centi`emes jusqu’`apresque une dizaine d’unit´es astronomiques, orbites circulaires ou mˆeme tr`es ex- centriques. Et le nombre de d´ecouvertes n’arrˆete pas d’augmenter.

Motivation, contexte, et objectifs

Les plan`etes g´eantes orbitant tr`es proche de leurs ´etoiles ont ´et´eappel´ees Jupiters chauds. Cette proximit´eaugmente les chances d’occurrences (voir Equation´ 1) d’un ph´eno- m`ene tr`es important, celui d’un plan´etaire, quand la plan`ete croise le disque de son ´etoile du point de vue d’un observateur sur Terre. Les transits plan´etaires repr´esentent une m´ethode compl´ementaire `acelle de vitesse radiale Doppler, ce qui permet une description plus compl`ete des propri´et´es physiques plan´etaires et stellaires. En combinant des mesures de vitesse radiales avec des mod`eles qui d´ecrivent la courbe de lumi`ere obtenue pendant un transit, il est possible de d´eterminer l’angle d’inclinaison de l’orbite (et par cons´equent la vraie masse de la plan`ete au lieu d’une limite inf´erieure simplement), le rapport entre les rayons de la plan`ete et de l’´etoile, et si on connaˆıt le rayon stellaire on peut aussi d´eterminer le rayon plan´etaire, sa densit´emoyenne et la gravit´e`ala surface. Les tran- sits d’une plan`ete en face d’une ´etoile brillante permettent, en plus, une d´etermination pr´ecise du rayon plan´etaire et de sa densit´emoyenne, ce qui est dˆu`ala haute qualit´edes mesures photom´etriques disponibles. Spectres `ahaut rapport signal-sur-bruit et haute r´esolution peuvent aussi ˆetre obtenus, et une ´etude d´etaill´ee de la composition chimique du syst`eme peut ˆetre effectu´ee. Toutes ces informations pourraient mettre des contraintes assez importantes dans les mod`eles th´eoriques qui d´ecrivent la structure et la constitution des plan`etes extrasolaires.

i R´esum´een fran¸cais

En plus de la courbe de lumi`ere, les transits peuvent aussietre ˆ observ´es `apartir d’un effet appel´eRossiter-McLaughlin, selon lequel les vitesses radiales sont d´evi´ees en plus de la d´eviation caus´ee par la pr´esence d’une plan`ete. La mod´elisation de ce ph´enom`ene nous permet de v´erifier `aquel niveau l’´equateur stellaire est align´eavec le plan de l’orbite de la plan`ete. Cette ´etude, qui est possible seulement pour des ´etoiles suffisamment brillantes pour permettre des mesures pr´ecises de vitesse radiale, peut nous aider `amieux comprendre les m´ecanismes de formation et d’´evolution des exoplan`etes.

Jusqu’`apr´esent, 30 plan`etes avec transits ont ´et´ed´etect´ees, mais seulement quelques unes orbitent des ´etoiles suffisamment brillantes pour permettre une ´etude d´etaill´ee. Malgr´e le progr`es concernant notre connaissance des m´ecanismes de formation et ´evolution des exoplan`etes, nous sommes encore assez loin d’avoir une description compl`ete de leurs propri´et´es physiques principales. La d´ecouverte et caract´erisation de nouveaux Jupiters chauds pourrait sˆurement nous aider `ar´epondre `ades questions encore ouvertes.

La strat´egie de ce travail de th`ese est essentiellement bas´ee sur deux aspects. Le premier vient de certaines ´etudes statistiques de l’´echantillon de plan`etes connues, ce qui a mis en ´evidence quelques effets syst´ematiques tr`es int´eressants. Il a ´et´ed´emontr´eque les ´etoiles ayant des plan`etes g´eantes en orbites tr`es proches ont un contenu en m´etal qui est, en moyenne, significativement plus grand (par environ 0.25 dex dans l’´echelle de m´etallicit´e[Fe/H]) que les ´etoiles isol´ees autour desquelles aucune plan`ete n’a ´et´etrouv´ee. Autrement dit, les ´etoiles riches en m´etaux (surtout celles avec [Fe/H] ! 0.2 dex) ont une plus grande probabilit´ed’abriter des plan`etes g´eantes proches que les ´etoiles ayant le mˆeme contenu en m´etal (ou moins) que le Soleil. L’existence de ce lien entre la pr´esence d’une plan`ete et la haute m´etallicit´ede l’´etoile hˆote a ´et´e propos´ee pour la premi`ere fois par Gonzalez (1997, 1998), et confirm´ee ult´erieurement par plusieurs travaux. Dans l’analyse statistique r´ealis´ee par Santos et al. (2001, 2004) et Fischer & Valenti (2005), les auteurs ont pu v´erifier qu’environ seulement 3% des ´etoiles de m´etallicit´esolaire ont des plan`etes d´etect´ees, tandis qu’environ 25% des ´etoiles avec [Fe/H] > 0.3 dex abritent un compagnon plan´etaire (voir Figure 1).

Le deuxi`eme aspect de notre strat´egie de travail ´etait de focaliser les observations dans la recherche de plan`etes `acourte p´eriode. Pour cela, chacune des ´etoiles riches en m´etaux ´etait normalement observ´ee deux jours apr`es la premi`ere mesure afin de pou- voir ´eventuellement identifier des variations `acourte p´eriode dans leurs vitesses radiales. De cette fa¸con, on esp´erait pouvoir augmenter nos chances de d´ecouvrir des nouveaux syst`emes extrasolaires pour lesquels des transits plan´etaires ont une plus grande proba- bilit´ed’occurrence. La probabilit´ed’occurrence d’un transit est donn´ee par la relation suivante, de Bodenheimer et al. (2003):

1 AU R +R 1 e cos (π/2 ω) (transit) = 0.004649 ! p − − (1) P a R! 1 e2 ! "! "# − $ o`u R! et Rp sont, respectivement, le rayon stellaire et plan´etaire (en unit´es de rayon solaire), a le demi-grand axe (en UA), e l’excentricit´e, et ω la longitude du p´eriastre.

ii S´election de l’´echantillon et observations

Figure 1: En haut, `agauche: Distribution de m´etallicit´epour une liste d’´etoiles `a plan`ete (histogramme hachur´e) et pour des ´etoiles de comparaison sans plan`etes d´etect´ees (his- togramme ouvert). La diff´erence moyenne de m´etallicit´eentre les deux ´echantillons est de 0.25 dex. En haut, `adroite: Fonctions cumulatives des deux distributions. En bas, `a gauche: Distribution de m´etallicit´epour des ´etoiles avec plan`etes inclues dans l’´echantillon CORALIE de recherche de plan`etes (histogramme hachur´e) et pour 875 ´etoiles dans l’ensemble du programme CORALIE ayant au moins 5 mesures (histogramme ouvert). En bas, `adroite: Pourcentage d’´etoiles avec plan`etes trouv´ees dans l’´echantillon CORALIE par bin de m´etallicit´e. Figure tir´ee de Santos et al. (2004).

S´election de l’´echantillon et observations

C’est en tenant compte de ces r´esultats des ´etudes statistiques mentionn´ees ci-dessus que l’on a cr´e´eun programme d’observation bas´esur un ´echantillon d’´etoiles dont seule- ment celles avec un haut contenu en m´etal ont ´et´esuivies. L’´echantillon inclut des ´etoiles de type solaire du catalogue Hipparcos (ESA 1997) s´electionn´ees de la fa¸con suivante:

iii R´esum´een fran¸cais

i) ´etoiles dans l’h´emisph`ere nord (d´eclinaison δ 0), avec une distance limit´ee `a100 pc, ≥ − et suffisamment brillantes (V 8.5) pour permettre une bonne pr´ecision ( 15 m s 1 ≤ ∼ ou moins) dans les mesures individuelles de vitesse radiale. ii) F8 type spectral < M1, ce qui correspond `a0.45 < (B V ) 1.4. Les ´etoiles ≤ − ≤ plus chaudes que le type solaire ont un nombre plus petit de raies spectrales, et des raies avec des profils plus larges. Ces effets diminuent la pr´ecision dans les mesures de vitesses radiales. D’un autre cot´e, les ´etoiles plus froides sont trop faibles pour garantir une bonne pr´ecision en vitesse radiale avec le spectrographe ELODIE. iii) on a ´elimin´eles ´etoiles ´evolu´ees, en s´electionnant seulement celles ´etant au maximum 1.5 mag au-dessus de la s´equence principale. Les ´etoilesevolu´ees ´ sont connues pour avoir des variations intrins`eques en vitesse radiale caus´ees e.g. par des pulsations. iv) on a aussi exclu les ´etoiles d´ej`apr´esentes dans d’autres programmes de recherche de plan`etes, comme les anciens programmes ELODIE ou les programmes li´es au t´elescope Keck. v) finalement, en utilisant le catalogue Hipparcos et la base de donn´ees du spectrographe CORAVEL (Benz & Mayor 1984; Duquennoy et al. 1991; Duquennoy & Mayor 1991), on a identifi´eet ´elimin´eles ´etoiles binaires et les ´etoiles ayant une haute vitesse de rotation. Cela inclut: binaires spectroscopiques `adeux spectres; binaires serr´ees avec une s´eparation angulaire inf´erieure `a10 secondes d’arc; ´etoiles avec une erreur en vitesse radiale plus grand que 2 km s−1; et ´etoiles avec v sin i>6 km s−1. Les ´etoiles avec une haute vitesse de rotation ont des raies spectrales ´elargies, ce qui augmente l’incertitude de la vitesse radiale calcul´ee. En plus, la vitesse de rotation est directement li´ee au niveau d’activit´echromosph´erique des ´etoiles de la s´equence principale (Noyes et al. 1984), dans le sens o`ules ´etoiles tournant plus vite sont plus actives. L’occurrence de certains ph´enom`enes physiques dans la chromosph`ere des ´etoiles actives, comme la pr´esence de taches, peut perturber ou mˆeme simuler les variations en vitesse radiale caus´ees par la pr´esence d’une plan`ete.

Un catalogue contenant un total de 1061 ´etoiles de type solaire a donc ´et´ecr´e´e. La Figure 2 montre le diagramme HR (`agauche) et le diagramme couleur-distance (`adroite) pour les ´etoiles s´electionn´ees.

Notre programme d’observations a d´ebut´een Mars 2004 `al’Observatoire de Haute- Provence (France) en utilisant le spectrographe ´echelle ELODIE (Baranne et al. 1996), mont´eau foyer du t´elescope de 1.93 m. Cet instrument a une r´esolution spectrale (λ/∆λ) de 42 000. Avec des poses typiques de 20 min, les spectres obtenus ont un rapport signal- sur-bruit entre 20 et 100, ce qui correspond `aun bruit de photon entre 5 et 20 m s−1 ∼ ∼ dans les mesures individuelles. Un total de 2341 spectres de 854 ´etoiles a ´et´eobtenu pendant la dur´ee du programme, qui ´etait pr´evue pour au moins 3 ans.

La strat´egie des observations ´etait la suivante: apr`es la premi`ere mesure de chaque ´etoile, on utilisait la Fonction de Corr´elation Crois´ee du spectre stellaire pour obtenir

iv S´election de l’´echantillon et observations

Figure 2: Diagramme HR (gauche) et diagramme couleur-distance (droite) pour les 1061 ´etoiles s´electionn´ees. La ligne `atiret dans le diagramme de gauche indique la limite en magnitude choisie pour ´eliminer les ´etoiles ´evolu´ees.

une estimation de la m´etallicit´e[Fe/H] et de la vitesse de rotation projet´ee v sin i. La m´etallicit´eest bas´ee sur une calibration de [Fe/H] en fonction de la surface de la fonction de corr´elation et sur l’index de couleur B V de l’´etoile (similaire `acelle utilis´ee par − Naef 2004). La vitesse de rotation projet´ee provient d’une calibration de la largeur de la fonction de corr´elation en fonction de B V (Queloz et al. 1998). Les ´etoiles ayant [Fe/H] − − 0.1 dex et v sin i 6 km s 1ont ´et´es´election´ees afin d’ˆetre suivies pendant le programme. ≥ ≤ La Figure 3 donne une vue globale de certains param`etres li´es `ades mesures r´ealis´ees avec ELODIE: signal-sur-bruit, bruit photonique, v sin i, [Fe/H], nombre de mesures par ´etoile et dispersion en vitesse radiale.

Le spectrographe ELODIE a ´et´ed´ecommission´een Aoˆut 2006, avant la conclusion de notre programme. Il a ´et´eremplac´epar un instrument beaucoup plus efficace, le spectro- graphe SOPHIE (Bouchy et al. 2006), qui peut atteindre une pr´ecision de 3-4 m s−1 dans les mesures de vitesse radiale. Un consortium Franco-Suisse (Loeillet et al. 2006), men´e par Fran¸cois Bouchy, a ´et´ecr´e´eavec l’objectif d’explorer les capacit´es du nouvel instrument dans un large programme scientifique, compos´ede cinq sous-programmes. Cet instrument pourra continuer `aobserver 1) les ´etoiles sans aucune mesure faite avec ELODIE, 2) les ´etoiles avec seulement 1 ou 2 mesures et ayant [Fe/H] 0.1, et 3) les ´etoiles ayant une ≥− variation importante en vitesse radiale mais pour lesquelles aucune solution orbitale n’a ´et´etrouv´ee avec les mesures ELODIE. Un total de plus de 400 ´etoiles satisfait tous ces crit`eres et a ´et´edonc pris en compte par le consortium.

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Figure 3: Vue globale des observations montrant: la distribution du signal-sur-bruit (S/B) des spectres obtenus et le bruit de photon en fonction de S/B (en haut); les distributions de la vitesse de rotation projet´ee et de la m´etallicit´e(au milieu); et les distributions du nombre de mesures effectu´ees pour chaque ´etoile ayant [Fe/H] 0.1 et v sin i 6 km s−1, ≥ ≤ et de la dispersion en vitesse radiale des ´etoiles avec au moins 3 mesures (en bas).

vi R´esultats et discussion

Table 1: R´esum´edes principaux param`etres orbitaux des syst`emes plan´etaires d´ecouverts ou confirm´es avec le pr´esent travail de th`ese. Ceux annonc´es par d’autres groupes sont ´egalement list´es `ala fin.

Nom de la P Mp sin i K a e −1 [Fe/H]! R´ef. Plan`ete [jours] [MJup] [m s ] [UA] HD 189733 b 2.2190 ± 0.0005 1.15 0.0 (fix´ee) 205 ± 6 0.0313 −0.03 1 HD 219828 b 3.8335 ± 0.0013 0.06 0.0 (fix´ee) 7.0 ± 2 0.048 0.19 9 HD 149143 b 4.088 ± 0.006 1.36 0.08 ± 0.04 163 ± 8 0.052 0.20 6, 3 HD 102195 b 4.1138 ± 0.0006 0.45 0.0 (fix´ee) 63 ± 0.5 0.05 0.05 5, 9 HD 118203 b 6.1335 ± 0.0006 2.13 0.309 ± 0.014 217 ± 3 0.07 0.15 3 HD 185269 b 6.8399 ± 0.0013 1.03 0.225 ± 0.025 93.6 ± 2.5 0.077 0.10 10, 8 HD 43691 b 36.96 ± 0.02 2.49 0.14 ± 0.02 124 ± 2 0.24 0.28 4 HD 45652 b 43.67 ± 0.08 0.60 0.25 ± 0.04 37.3 ± 1.6 0.24 0.15 12 HD 132406 b 974 ± 39 5.61 0.34 ± 0.09 115 ± 26 1.98 0.18 4

HD 17156 b 21.2 ± 0.3 3.12 0.67 ± 0.08 275 ± 15 0.15 0.24 7 HD 155358 b 195.0 ± 1.1 0.89 0.112 ± 0.037 34.6 ± 3.0 0.63 −0.68 2 HD 75898 b 418.2 ± 5.7 2.51 0.10 ± 0.05 58.2 ± 3.1 1.19 0.27 11 HD 155358 c 530.3 ± 27.2 0.5 0.176 ± 0.174 14.1 ± 1.6 1.22 −0.68 2

1 Bouchy et al. (2005), 2 Cochran et al. (2007), 3 Da Silva et al. (2006), 4 Da Silva et al. (2007), 5 Ge et al. (2006), 6 Fischer et al. (2006), 7 Fischer et al. (2007), 8 Johnson et al. (2006), 9 Melo et al. (2007), 10 Moutou et al. (2006), 11 Robinson et al. (2007b), 12 Santos et al. (2008)

R´esultats et discussion

A` pr´esent, 13 plan`etes extrasolaires ont ´et´etrouv´ees en orbite autour d’´etoiles faisant partie de notre ´echantillon principal. Six ont ´et´ed´ecouvertes et annonc´ees par notre groupe, deux ont ´et´eult´erieurement confirm´ees par nos mesures, et quatre ont ´et´ean- nonc´ees par d’autres groupes de recherches des exoplan`etes. La derni`ere de nos d´ecouvertes a ´et´er´ecemment d´etect´ee par nos mesures de vitesses radiales et le papier em question est en pr´eparation. La Table 1 liste les principales propri´et´es de ces 13 exoplan`etes, ainsi que la m´etallicit´ede leurs ´etoiles et les papiers de publication respectifs. Il est important de noter que la plan`ete en orbite autour de l’´etoile HD 219828 est dans le r´egime de plan`etes de basse masse (Mp sin i = 19.8M⊕), n’´etant pas d´etectable avec la pr´ecision de notre programme. Les mesures de vitesse radiale de cette ´etoile ont ´et´efaites avec le spectro- graphe HARPS, en utilisant un sous-´echantillon de notre programme form´ed’´etoiles `a basse d´eclinaison, dans le cadre d’un programme d’observation men´epar C.H.F Melo du D´epartement d’Astronomie de l’Universit´edu Chili.

Parmi les plan`etes d´etect´ees, six ont des p´eriodes orbitales plus courtes que 10 jours, ayant donc une plus grande probabilit´ede pr´esenter un transit. En effet, le r´esultat le plus important de ce travail est li´e`al’´etoile HD 189733, pour laquelle le ph´enom`ene de transit plan´etaire a ´et´epremi`erement d´etect´epar nos observations (Bouchy et al. 2005).

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Table 2: Quelques param`etres physiques concernant le syst`eme HD 189733, par ordre de 1 l’ann´ee de publication. Toutes les r´ef´erences list´ees adoptent M =0.82 0.03 M! pour la ! ± masse de l’´etoile. J’ai recalcul´eles valeurs de densit´e en utilisant 1 R = 7.1492 107 m Jup × et 1 M = 1.8986 1027 kg. Les chiffres entre parenth`eses indiquent la r´ef´erence adopt´ee Jup × pour le param`etre respectif.

Rp R! Mp ρp i Tp R´ef. −3 [RJup] [R"] [MJup] [kg m ] [deg] [K] 1 1.26 ± 0.03 0.76 ± 0.01 1.15 ± 0.04 710 85.3 ± 0.1 - 2 (1) (1) (1) 710 - 1117 ± 42b 3 1.154 ± 0.033 0.758 ± 0.016a (1) 930 85.79 ± 0.24 1279 ± 90b 4 1.10 ± 0.03 0.73 ± 0.02 1.13 ± 0.03 1050 86.1 ± 0.2 - 5 1.156 ± 0.046 0.753 ± 0.025 (4) 910 85.76 ± 0.29 - 6 1.19 ± 0.08 0.779 ± 0.052 (1) 850 - - 7 1.137 ± 0.006 0.757 ± 0.003 (1) 970 85.61 ± 0.04 1205.1 ± 9.3c 8 1.154 ± 0.017 0.755 ± 0.011 (1) 930 85.68 ± 0.04 -

1 Bouchy et al. (2005), 2 Deming et al. (2006), 3 Bakos et al. (2006b), 4 Winn et al. (2006), 5 Winn et al. (2007a), 6 Baines et al. (2007), 7 Knutson et al. (2007b), 8 Pont et al. (2007b) a publi´epar Masana et al. (2006), b temp´erature de brillance `a16 µm, c temp´erature de brillance `a8 µm

Depuis son annonce, plusieurs travaux concernant ce syst`eme ont ´et´epubli´es (voir Table 2), en faisant une ´etude d´etaill´ee de ses propri´et´es physiques et g´eom´etriques, ainsi qu’une description de la composition chimique de l’atmosph`ere de la plan`ete. Notamment, la proximit´edu syst`eme (19 pc) et le haut rapport entre les rayons de la plan`ete et de l’´etoile (> 0.15) offrent des conditions tr`es favorables `ades ´etudes photom´etriques et spectroscopiques pr´ecises.

Un transit plan´etaire concernant le syst`eme HD 17156 a ´et´er´ecemment publi´epar Barbieri et al. (2007). La plan`ete dans ce syst`eme a ´et´epremi`erement annonc´ee par Fischer et al. (2007) en utilisant la technique de vitesse radiale, et son mouvement orbital a pu aussi ˆetre caract´eris´eavec nos mesures de vitesse radiale. La particularit´ede ce syst`eme, une plan`ete avec une grande p´eriode orbitale (21 jours) et une grande excentricit´e(0.67) en comparaison avec les autres syst`emes `atransit connus, et le fait que l’´etoile principale soit assez brillante (V =8.17), nous donneront sˆurement des r´esultats tr`es int´eressants dans le futur proche, comme cela a ´et´ele cas de la HD 189733 et d’autres syst`emes `atransit ayant une ´etoile brillante.

Quelques unes des ´etoiles observ´ees dans le cadre de ce programme ont pr´esent´eune grosse variation en vitesse radiale, se r´ev´elant ˆetre membres d’un syst`eme binaire, ou avoir un compagnon faible non connu, comme une naine brune par exemple. Certaines de ces ´etoiles ont aussi des mesures effectu´ees avec le spectrographe CORAVEL, et la combinaison de donn´ees disponibles a permis d’ajuster une solution orbitale pour 4 ´etoiles: HD 62923 a probablement une naine brune comme compagnon, un objet de M2 sin i =0.051 M! orbitant l’´etoile primaire avec une p´eriode de 124 jours. HD 89010 a comme compagnon

viii Analyse statistique du programme

une ´etoile avec M2 sin i =0.274 M! en orbite de p´eriode 788 jours. L’objet autour de HD 212733 a une masse M2 sin i =0.109 M! et une p´eriode P = 90 jours. Et le dernier objet a une masse M2 sin i =0.109 M! et orbite autour de l’´etoile HD 216320 avec une p´eriode de 20 jours.

Analyse statistique du programme

Afin de d´eterminer quelles sont les limites de d´etection dans notre recherche de Jupiters chauds en utilisant le spectrographe ELODIE, j’ai effectu´e des simulations num´e-riques Monte Carlo bas´ees dans l’ensemble de donn´ees collect´ees pendant la dur´ee du programme. Ces simulations sont capables d’estimer quel est le niveau de d´etection pour une certaine masse plan´etaire et une certaine p´eriode orbitale donn´ees, en tenant compte de la masse de l’´etoile, du calendrier de mesures de vitesse radiale, et de l’incertitude dans chaque mesure individuelle. Pour chaque valeur de masse et de p´eriode orbitale, j’ai cr´e´e1000 orbites Kepleriennes al´eatoires, et pour chaque orbite les param`etres suivants ont ´et´echoisis de la fa¸con d´ecrite:

Etoile´ de r´ef´erence: choisie al´eatoirement d’une liste de 166 ´etoiles riches en m´etaux ([Fe/H] 0.1), de basse rotation (v sin i 6 km s−1), et ayant au moins 3 mesures ≥ ≤ de vitesse radiale. Pour chaque ´etoile s´electionn´ee, sa masse, la date Julienne des mesures et les bruits de photons respectifs sont pris en compte.

Erreur instrumentale et bruit stellaire intrins`eque: de la liste de 166 ´etoiles cr´e´e ci-dessus, j’ai s´electionn´eseulement celles pour lesquelles aucun compagnon stellaire ou plan´etaire n’a ´et´ed´etect´e, et pour chacune j’ai calcul´ela dispersion en vitesse ra- diale corrig´ee du bruit photonique moyen des mesures. Pendant les simulations, une valeur de cette dispersion, qui repr´esente les sources d’erreurs non-photoniques, est tir´ee au hasard et somm´ee quadratiquement au bruit photonique de l’´etoile choisie.

Instant de passage au p´eriastre (T0): tir´eau hasard uniform´ement dans l’intervalle t1 et t1 + P , o`u t1 repr´esente la date Julienne de la premi`ere mesure de vitesse radiale de l’´etoile choisie, et P est la p´eriode orbitale donn´ee.

Excentricit´eorbitale (e): pendant les simulations, j’ai fix´el’excentricit´e`az´ero, ce qui est une bonne approximation ´etant donn´eque les plan`etes que l’on a cherch´ees sont principalement des Jupiter chauds, qui ont en g´en´eral des orbites circulaires. En tous cas, j’ai compar´eles r´esultats avec e = 0 aux r´esultats en fixant l’excentricit´e`a e =0.5 et e =0.8, et les diff´erences dans la probabilit´ede d´etection sont de l’ordre de 10% ou moins.

Angle d’inclinaison orbitale (i): des valeurs de sin i ont ´et´etir´ees au hasard et uni- form´ement dans l’intervalle 0 et 1.

ix R´esum´een fran¸cais

Figure 4: Probabilit´ede d´etection de variations en vitesse radiale en fonction de la p´eriode orbitale pour plusieurs valeurs de masse stellaire, de 0.3a ` 10 MJup. Notez la diminution de la probabilit´epour des p´eriodes `a1, 2 et 365 jours.

Dans l’´etape suivante, des valeurs simul´ees de vitesse radiale sont calcul´ees aux mo- ments donn´es par les dates Juliennes, et en tenant en compte de l’erreur respective de chaque mesure. Finalement, les dispersions des vitesses simul´ees sont calcul´ees. Si ces dispersions d´epassent une certaine limite pr´e-d´efinie, alors j’ai consider´eque les variations en vitesse radiale sont d´etectables. Sinon, j’ai consider´eque les possibles variations ne peu- vent pas ˆetre d´etect´ees, et que l’´etoile est donc constante. La Figure 4 montre les r´esultats des simulations avec un plot de la probabilit´ede d´etection d’une variation consid´erable en vitesse radiale en fonction de la p´eriode orbitale de la plan`ete, pour diff´erentes valeurs de masse stellaire.

J’ai utilis´eles simulations pour d´eterminer la probabilit´ede d´etection de plan`etes g´eantes dans les points (P,Mp) correspondant `achaque d´etection d’´etoiles avec plan`etes pr´esentes dans notre programme. Ensuite, j’ai utilis´eces probabilit´es pour estimer quel est le nombre de plan`etes qui pourraient ˆetre effectivement trouv´ees si aucun biais obser-

x Analyse statistique du programme

Table 3: Estimation du nombre de d´etections (Npl) corrig´edes sources d’erreurs non- photoniques. Le param`etre est la probabilit´ede d´etection d’une certaine plan`ete Pdet donn´ees sa masse et sa p´eriode orbitale (N =1/ ). pl Pdet P<10 jours:

Mp sin i P det Npl [MJup] [jours] P HD 149143 b 1.36 4.09 0.66 1.52 HD 118203 b 2.13 6.13 0.74 1.35 HD 185269 b 1.03 6.84 0.50 2.00 subtotal 5 2 ∼ ± 10 P<100 jours: ≤ HD 17156 b 3.12 21.20 0.68 1.47 HD 43691 b 2.49 36.96 0.54 1.85 HD 45652 b 0.60 43.70 0.08 12.5 subtotal 16 4 ∼ ± 100 P<1000 jours: ≤ HD 75898 b 2.51 418.20 0.14 7.14 HD 132406 b 5.61 974.00 0.26 3.85 subtotal 11 3 ∼ ± Total 32 5 ∼ ± vationel n’´etait pr´esent. La Table 3 liste les r´esultats. On peut v´erifier que le nombre de plan`etes d´ej`atrouv´ees par ce programme n’est pas loin de ce que l’on peut attendre si l’on consid`ere seulement les cas pour lesquels la probabilit´e de d´etection est d’au moins 50%. La Table 1 liste un total de 12 ´etoiles avec plan`ete. Deux ont des plan`etes avec une p´eriode orbitale plus longue que 100 jours et, en mˆeme temps, sont d´etectables avec ELODIE. On peut esp´erer trouver encore plusieurs autres, mais cela d´epend de la continuation de notre programme qui sera, en partie, inclus dans les programmes de recherche de plan`etes du nouveau spectrographe SOPHIE. La haute pr´ecision en vitesse radiale atteinte par cet instrument (3-4 m s−1) en comparaison avec ELODIE (10-15 m s−1) va probablement con- tribuer `acompl´eter la liste d’´etoiles avec plan`etes qui sont probablement pr´esentes dans notre ´echantillon.

xi R´esum´een fran¸cais

xii Contents

R´esum´een fran¸cais i Motivation, contexte, et objectifs ...... i S´election de l’´echantillon et observations ...... iii R´esultatsetdiscussion...... vii Analyse statistique du programme ...... ix

1 Introduction 1

I The search for short-period planets in metal-rich stars 5

2 Motivation, context and objectives 7 2.1 The metallicity of stars hosting planets ...... 7 2.2 The planet search biased towards metal-rich stars ...... 10 2.3 Planetary transits and the search for hot Jupiters ...... 10 2.3.1 Planets transiting bright stars ...... 12 2.3.2 Transmissionspectroscopy...... 16 2.3.3 Occultation photometry and spectroscopy ...... 17 2.3.4 The Rossiter-McLaughlin effect...... 19 2.3.5 The transit light curve ...... 20 2.4 Other sources of RV variability ...... 24 2.4.1 The stellar chromospheric activity ...... 25 2.4.2 Line bisector analysis ...... 27

3 Sample selection and observations 31 3.1 Selectionofthetargetstars ...... 31 3.2 Observations ...... 33 3.3 Calibrations using the CCF profile ...... 35 3.3.1 Calibration of the projected rotation velocity ...... 36 3.3.2 Calibration of the metal abundance ...... 36 3.3.3 ThedependenceonS/N ...... 38

xiii Contents

3.4 Global view of the observations ...... 40 3.4.1 Photon-noiseerrors...... 42

4 Main results 51 4.1 The discovery of a transiting HD 189733 ...... 51 4.1.1 Paper: A very hot Jupiter transiting the bright star HD189733 . . . 52 4.2 Other publications concerning the HD 189733 system ...... 59 4.3 Other hot Jupiters discovered with our programme ...... 62 4.3.1 Paper: Hot Jupiters around HD 118203 and HD 149143 ...... 63 4.3.2 Paper: A hot Jupiter orbiting HD 185269 ...... 71 4.4 Additional discoveries ...... 77 4.4.1 Paper: Two planets orbiting HD 43691 and HD 132406 ...... 83 4.4.2 Paper: Two hot companions around HD 102195 and HD 219828 . . . 91 4.5 Radial-velocity variable stars with no detected planet ...... 101 4.5.1 Photometric variable and/or active stars ...... 101 4.5.2 Somepeculiarcases ...... 106 4.5.3 Binaryorbrowndwarfstars...... 111

5 Statistical study of the programme 117 5.1 Metallicity distribution ...... 117 5.2 Period distribution ...... 119 5.3 Global detectability of the programme ...... 119 5.3.1 Simulations: description ...... 121 5.3.2 Detection limits in the planetary ...... 130 5.3.3 Simulations: results ...... 131

6 Conclusion and prospects 133 6.1 Summary of our results and main conclusions ...... 133 6.2 Newperspectives ...... 135

II Appendices 137

A Sample stars 139

xiv Chapter 1

Introduction

After the announcement of 51 Peg b by Mayor & Queloz (1995), the first planet orbiting a star other than the Sun, many other authors have dedicated considerable efforts in the search for new worlds beyond the frontiers of our solar system. The unexpected physical properties of 51 Peg b, a large planet ( of 0.47 Jupiter masses) orbiting its parent stars with a short period (only 4.2 days), triggered even more motivation to detect such systems, with the aim of understanding the origin and nature of planets that are so different from what we know about those orbiting our Sun. To date, about 250 of these so-called extrasolar planets have been detected, with masses that range from some Earth masses to several Jupiter masses, orbital distances ranging from a few hundredths to almost 10 astronomical units, orbits that are circular or even very eccentric. And the number of candidates does not cease growing.

The close-in giant planets have been called hot Jupiters given their proximity to their parent star. A very special group of hot Jupiters is the one composed of planets that transit the disk of the main star as viewed from Earth. Transiting systems represent a complementary method of planet detection that, together with the Doppler radial-velocity technique, leads to a more complete description of the planetary and stellar properties. Combining radial-velocity observations with models that describe the transit light curve, we can derive the inclination angle of the orbit (which allows the estimate of the real mass of the planet instead of its lower limit), the radius ratio between star and planet, and if we know the stellar radius we can also determine the planetary radius, its mean density, and its surface gravity. In particular, transiting planets orbiting bright stars are of extreme importance, since the high quality of the photometric light curves attained for such systems provides a precise determination of the planetary radius and mean density. High signal-to-noise and high resolution spectra can also be gathered, making possible a detailed study of the chemical composition of the system. All this information can put significant constraints on models of the structure and constitution of extrasolar planets.

1 Chapter 1. Introduction

Besides the light curve, the transit event is also manifested by means of the Rossiter- McLaughlin effect, in which an anomalous radial-velocity shift is superimposed on the radial-velocity variation induced by the planetary motion. Modelling this phenomenon allows us to verify how closely the stellar equator is aligned with the orbital plane of the planet. Such a study, which is only possible for stars bright enough to provide accurate radial-velocity measurements, may help us to understand the mechanisms of planetary formation and evolution.

The primary motivation behind these studies comes from the anomalous results de- scribing the properties of transiting their host stars in comparison to what we know from the planets of our solar system. The close proximity of these objects to their stars has important consequences for their radius given the effects of tidal interactions and the intense stellar irradiation. In one hand, some of the discovered extrasolar tran- siting planets have radius larger than the expected by models that describe their internal structure. The strong irradiation from the star would delay the evolutionary contraction of the planet, but recent results have shown that, apparently, this effect is not enough to explain the observed radius. On the other hand, the opposite of this situation has also been observed, in which the planetary radius is smaller than the predicted by standard models of giant planet formation. Such a small radius indicates that the planet is made of a significant fraction of heavy elements, much more than what is normally assumed by the models.

To date 30 planets have been detected transiting their parent stars, but only a small fraction of them are bright enough to allow a detailed study. Despite many improvements in our knowledge concerning the mechanisms of formation and evolution of the exoplanets, it seems that we are still far from having a complete description of their main physical properties. The characterisation of new hot Jupiters would surely help to answer the main questions still open.

The strategy of the present work was based essentially on two aspects. The first of them comes from statistical studies of the current sample of known extrasolar planets, which put in evidence some interesting trends. In particular, it has been shown that stars harbouring giant planets in close orbits have a metal content that is, in average, significantly higher (by about 0.25 dex in the metallicity scale [Fe/H]) than single stars for which no planet was detected. In other words, metal-rich stars (specially those having [Fe/H] ! 0.2 dex) are more likely to host close-in giant planets than stars with solar metal content or poorer than the Sun. In the statistical analysis of Santos et al. (2001, 2004) and Fischer & Valenti (2005), the authors were able to verify that only about 3% of the stars with solar metallicity have detected planets, whereas roughly 25% of the stars with [Fe/H] > 0.3 dex were found to harbour a planetary companion.

For this reason, we created an observation programme based on a sample of 1061 stars from which we monitored only those presenting a high metal content. The whole sample includes low-rotation solar-type stars in the northern hemisphere, within a distance of

2 100 pc, and bright enough (V 8.5) to allow a good precision ( 15 m s−1 or smaller) in ≤ ∼ the individual radial-velocity measurements. In our observational strategy, after the first measurement of each star, we used the Cross Correlation Function (CCF) of the stellar spectrum to estimate a value for metallicity based on a calibration of [Fe/H] as a function of the CCF’s width and the star’s B V colour index. The stars having [Fe/H] 0.1 dex − ≥ were selected to be monitored during the duration of the programme, predicted to be of at least 3 years. The second aspect of our strategy was to focus our observations on the search for short-period planets: each one of the metal-rich stars was scheduled to be observed two days after the first measurement in order to eventually identify short-period variations in the radial velocities. Doing so, we expected to increase our chances to discover new extrasolar planets for which transit events across the disk of their parent stars are more likely to take place.

Up to now, 13 extrasolar planets have been discovered orbiting stars present in our main sample of 1061 objects. Six of them were first discovered and announced by our team, two were confirmed latter by our measurements, and four were announced by other groups of search for extrasolar planets. The last one was recently detected by our radial- velocity measurements, and the related paper has been submitted. Such announced new extrasolar planets are, in order of the first publication: HD189733 b (Bouchy et al. 2005); HD 118203 b (Da Silva et al. 2006); HD 149143 b (Fischer et al. 2006; Da Silva et al. 2006); HD 185269 b (Johnson et al. 2006; Moutou et al. 2006); HD 102195 b (Ge et al. 2006; Melo et al. 2007); HD 219828 b (Melo et al. 2007); HD 43691 b and HD 132406 b (Da Silva et al. 2007); HD 17156 b (Fischer et al. 2007); HD 155358 b and HD 155358 c (Cochran et al. 2007); HD 75898 b (Robinson et al. 2007b); and HD 45652 b (Santos et al. 2008). I notice that the planet orbiting the star HD 219828 is in the regime of low mass planets (Mp sin i = 19.8M⊕), and hence is not detectable with the precision of the present programme.

Six of these planets have periods shorter than 10 days, showing a higher probability of transit events. Indeed, our most important result concerns the star HD 189733, for which a transit event have been observed by Bouchy et al. (2005). Since then, many other works related to this system have been published, allowing a detailed study of its physical and geometrical parameters as well as some description of the chemical composition of the planet’s atmosphere. Notably, the proximity of the system (only 19 pc) and the high planet-to-star radius ratio (> 0.15) offer very special conditions for accurate photometric and spectroscopic studies.

Transits of the planet HD 17156 b have been recently observed by Barbieri et al. (2007). This planet was first announced by (Fischer et al. 2007) using the radial-velocity technique, and its orbital motion could also be characterised with our radial-velocity measurements. The peculiarity of this system, a planet with large (21 days) and large eccentricity (0.67) in comparison to the other known transiting systems, and the fact that the primary is a bright star (V =8.17), will bring interesting results in the near future, as it has been the case for HD 189733 and other transiting systems with bright stars.

3 Chapter 1. Introduction

In Chapter 2, I discuss the reasons for choosing to search for planets around metal- rich stars, and I situate the reader in the context of hot Jupiters and the importance of observing planetary transits. I also discuss other possible sources of radial-velocity variability, and the solutions used in this work to discriminate the various situations. In Chapter 3, I describe the sample of stars selected to be observed in the course of our programme, together with the criteria of selection and the observational strategy. In Chapter 4, I present the main results achieved with this work and the related publications. In Chapter 5, I made a global analysis of the programme results, comparing the number of discovered planets to the number of planets expected to be found with the adopted strategy. Finally, in Chapter 6, I conclude the present work and I discuss some perspectives and future directions.

4 Part I

The search for short-period planets in metal-rich stars

5

Chapter 2

Motivation, context and objectives

Contents

2.1 The metallicity of stars hosting planets ...... 7 2.2 The planet search biased towards metal-rich stars ...... 10 2.3 Planetary transits and the search for hot Jupiters ...... 10 2.4 OthersourcesofRVvariability ...... 24

2.1 The metallicity of stars hosting planets

Gonzalez (1997, 1998), by studying the planetary systems known at that time, noted that the parent stars in these systems are not typical in which concerns their chemical compositions: such stars have metallicities much higher than the average for nearby stars. He suggested for the first time the existence of a link between presence of giant planets in close orbits and the high metallicity of the host star, and that the search for radial-velocity variations should focus on metal-rich stars.

The first question that came out was: what is the nature of this link? Two models have been proposed to explain the observed bias: 1) primordial approach: the high stellar metallicity has a primordial origin, coming from the interstellar cloud from which the star was formed. In this model, the high metallicity would be a prerequisite for the formation of giant planets (see Podolak et al. 1993); 2) pollution approach: the observed high metallicity is related to the migration process (in which a gas giant is formed at a large distance from the star and then migrates inwards due to interactions with the circumstellar disk) and, during the migration, the material between star and planet would be accreted by the star,

7 Chapter 2. Motivation, context and objectives

which would have the envelope enriched in metals (see Lin et al. 1996). A third possibility would be a combination of both.

Santos et al. (2001) made, for the first time, a uniform and unbiased comparison between main-sequence stars with and without planets observed using the CORALIE spectrograph mounted on the 1.2-m Euler Swiss telescope at the La Silla Observatory (ESO, Chile). Their results indeed show that dwarf stars hosting giant planets in close orbits are more metal-rich than field stars with no detected planetary companions. They did not rule out the hypothesis of accretion of material on the stellar envelope during the migration process, but argued that the source of the high metallicity is most probably primordial. Santos et al. (2003) had an enlarged and uniform sample, with about 80 solar- type stars with planets, and their analysis confirmed previous conclusions. In particular, they investigated the dependence of the stellar metallicity on the mass of the convective envelope since, if the pollution model is correct, the mixing of material due to convection will be more efficient when the depth of the convective zone is larger, and the metallicity will be smaller. However, this was not observed in their sample, leading them to support the primordial approach.

Santos et al. (2004) made a statistic analysis of an even larger number of stars hosting planetary companions, a total of 98 stars, and confirmed the main idea proposed by previous works. They were able to quantify that solar-type stars with close-in giant planets are, on average, more metal-rich by 0.25 dex than field dwarf stars. Moreover, they observed that the frequency of planets is a strongly rising function of the stellar metallicity, at least when [Fe/H] > 0. The distribution seems to be flat for [Fe/H] < 0, and increase very steeply for higher metallicities up to [Fe/H] = 0.4: only 3% of the stars with solar metallicity have planets, whereas 25% of the stars with [Fe/H] > 0.3 were found to harbour a planetary companion (see Figure 2.1). The authors revised the spectroscopic analysis, improving the surface gravity determination, which were then in agreement with those derived by other works and with the trigonometric gravities derived using the parallaxes from the Hipparcos catalogue (ESA 1997). The new determinations did not affect the main conclusions proposed before. Santos et al. (2005) added 29 new planet hosts and 53 comparison stars, yielding an even larger and uniform sample with accurate spectroscopic parameters that reinforces the results presented in previous studies.

Fischer & Valenti (2005) presented their results from a high-resolution spectroscopic analysis for every FGK stars on the Keck, Lick, and Anglo-Australian Telescope planet search surveys. They found that less than 3% of the main-sequence stars with 0.5 < − [Fe/H] < 0.0 have giant planets in close orbits. The percentage rapidly increases for [Fe/H] above the solar metallicity: 25% of the stars with [Fe/H] > 0.3 have detected giant planets. They established an analytic relation for the probability of formation of gas giant planets with orbital periods shorter than 4 yr and radial-velocity amplitudes larger than 30 m s−1 as a function of the stellar metallicity: (planet) = 0.03 102[Fe/H]. Wyatt et P × al. (2007) attributed this relation to the protoplanetary disk mass distribution. Fischer & Valenti also presented many arguments in favour of the primordial origin of the enhanced

8 2.1. The metallicity of stars hosting planets

Figure 2.1: Upper panel, left: Metallicity distribution for stars with planets (hatched his- togram) and for a comparison list of stars without planets (open histogram). The average difference in metallicity between the two samples is about 0.25 dex. Upper panel, right: Cumulative functions for both distributions. Lower panel, left: Metallicity distribution for stars with planets included in the CORALIE planet-search sample (hatched histogram) and for 875 stars in the whole CORALIE programme having at least 5 measurements (open histogram). Lower panel, right: Percentage of stars with planets found in the CORALIE sample per metallicity bin. Figure from Santos et al. (2004). metallicity observed in stars with close-in giant planets. They argued that there is no evidence that accretion occurs in a different way if one compares stars with planets to those with no detected planet. Thus, host planet stars are more likely to be metal-rich throughout their interiors and not just only at their external layers.

Afterwards, Gonzalez (2006) made a review concerning the observed composition trends and their possible causes. The author concluded that the primordial hypothesis is indeed the best explanation for the high metallicity of stars with close-in giant planets, arguing that any evidence for self-enrichment among stars with planets was still weak.

9 Chapter 2. Motivation, context and objectives

He suggested that an orbital period bias, related to the migration process, might exist. A metallicity-dependent mechanism would affect the final orbital position of a migrating planet. This conclusion, however, needs more statistical support.

2.2 The planet search biased towards metal-rich stars

Taking into account the various results confirming the high-metallicity nature of main- sequence stars harbouring giant gas planets in close orbits and, by consequence, the higher probability to find such planets around metal-richer stars, some groups of search for ex- trasolar planets decided to focus their efforts on initiating specific research programmes in which only stars having high metal content were selected for radial-velocity follow-up.

The N2K (Next Two Thousand) project (Fischer et al. 2004, 2005) is a large consortium of American, Chilean, and Japanese astronomers, with a sample of 2000 closest, brightest, and metal-rich ([Fe/H] > 0.1) FGK stars to be observed using the Keck, Magellan, and Subaru telescopes. They expect the detection of about 60 new planets with minimal masses larger than 0.5 Jupiter masses and orbital periods shorter than 14 days, including a few transiting systems. See also Robinson et al. (2007a) and references therein for an overview of the N2K consortium and its main results.

The present work is part of another metallicity-biased search programme (Da Silva et al. 2006) started in March 2004 at the Haute-Provence Observatory, using the high-precision ELODIE fiber-fed echelle spectrograph on the 1.93-m telescope. We selected a sample of more than a thousand of non-evolved, non-binary, and slow-rotating solar-type stars in the northern hemisphere, at distances smaller than 100 pc and brighter than V =8.5. The strategy consisted of observing each target star at least once, estimating a value for metallicity, and monitoring only those with [Fe/H] 0.1. We expected to discover a few ≥ tens of new planets, from which a few being hot Jupiters with high probability of transiting their parent star. See Chapter 3 for a detailed description of the sample of target stars and the strategy of the observations.

2.3 Planetary transits and the search for hot Jupiters

The main goal of the present work is to search for planets in close-in orbits around their parent stars, with periods shorter than 10 days. These planets are normally called hot Jupiters given the high temperatures attained on their surfaces due to their proximity to the star. Such systems, in comparison to those with larger orbital periods, offer a higher probability of occurrence of a very interesting situation: when the planet passes in front of the stellar disk, providing a planetary transit.

10 2.3. Planetary transits and the search for hot Jupiters

Transits can only be detected if the planetary orbit is near the line-of-sight between the observer and the star. Taking into account all possible inclination angles of the orbit, the probability that a planet crosses the stellar disk as viewed from Earth can be estimated using the following equation, from Bodenheimer et al. (2003):

1 AU R +R 1 e cos (π/2 ω) (transit) = 0.004649 ! p − − (2.1) P a R! 1 e2 ! "! "# − $ where R! and Rp are, respectively, the stellar and the planetary radius (both in units of solar radius), a is the orbital semimajor axis (in AU), e the eccentricity, and ω the argument of periastron.

The phenomenon of a planet transiting its host star as viewed from Earth allows the evaluation of some physical parameters of the system other than those we normally ob- tain by only doing radial-velocity measurements in non-transiting systems. Combining radial-velocity observations with models that describe the transit light curve (see Subsec- tion 2.3.5), one can estimate: 1) the exact value of the inclination angle (i) of the orbit, removing the sin i ambiguity; 2) the real mass of the planet (Mp) instead of the minimal value assigned by Mp sin i; 3) and the radius ratio between star and planet. By conse- quence, if the stellar radius is known, one can derive the planetary radius, its mean density (ρp), and its surface gravity (g). According to current models, intermediate-mass planets comprise some or all of the following four layers: iron/nickel core, silicates (rocks), ice (H2O, CH4, NH3), and an H/He envelope. The position of a planet in the mass-radius di- agram may give an indication of its constitution since, for a given mass, age, and distance to the parent star, the radius of the planet mostly depends on its composition. Therefore, a precise determination of the planet’s radius and mean density would shed light on our understanding of the structure and constitution of the known extrasolar planets.

Another characteristic of a transit event is that the radial-velocity measurements are affect by a phenomenon called Rossiter-McLaughlin effect, which allows us to verify how closely the stellar equator is aligned with the orbital plane of the planet, and may help us to understand the mechanisms of planetary formation and evolution (see more details in Subsection 2.3.4).

To date 30 planets transiting their parent stars have been detected and confirmed by both photometric and radial-velocity measurements. Table 2.1 lists some of the main properties of these systems1, which were derived thanks to the combined information provided by spectroscopic (from the radial-velocity measurements) and photometric (from the transit event) observations. In this table, some of the planets are named according to their discovery project: HAT is the Hungarian Automated Telescope project (Bakos et al. 2002, 2004); TrES is the Trans-Atlantic Survey (Alonso et al. 2004); OGLE is the Optical Gravitational Lensing Experiment (Udalski et al. 1992); XO is an exoplanet search programme (McCullough et al. 2005); and WASP is the Wide Angle

1 A complete list of the system properties is available in The Extrasolar Planets Encyclopaedia (http://exoplanet.eu) and references therein.

11 Chapter 2. Motivation, context and objectives

HD209458b

Figure 2.2: Evolution on the radius of a planet with mass Mp =0.69 MJup for different models: non-irradiated (long-dashed line), only irradiated (solid line), and irradiated with an extra source of energy deposited in layers between the surface and the mass shell M ∆m (the other three curves). The position of HD 209458 b is also indicated. Figure p − from Baraffe et al. (2003).

Search for Planets consortium (Pollacco et al. 2006). For comparison, Table 2.2 shows the radius, mass, and density of the giant planets of our solar system. Table 2.3 lists the main properties of the 30 host stars. Below, I give a special attention to a particular subset of transiting systems: those with planets transiting bright stars.

2.3.1 Planets transiting bright stars

The assortment of physical parameters derived from transiting systems can be even more enlarged when a planet transits a bright star (V<12). Such a situation yields high resolution and high signal-to-noise spectra, as well as high quality photometry, allowing a precise determination of the planetary parameters. HD 209458 b is the first extraso- lar planet known to transit its host star (Charbonneau et al. 2000; Henry et al. 2000; Mazeh et al. 2000). This represents the first time that a planet detected by the radial- velocity method could be confirmed by a completely independent technique. Together with HD 149026 b (Sato et al. 2005), HD 189733 b (Bouchy et al. 2005), HAT-P-2 b (Bakos et al. 2007b), and HD 17156 b (Barbieri et al. 2007) they form a singular subset of five systems with planets transiting stars brighter than V = 9.

A pertinent motivation to perform a detailed study of such transiting systems is related to some anomalous results describing their properties when compared to those we know

12 2.3. Planetary transits and the search for hot Jupiters

Table 2.1: The 30 transiting planets known to date. Those marked in bold represent the planets orbiting stars brighter than V = 9. The mean densities were calculated assuming the values below for the equatorial radius and mass of Jupiter. i P Rp Mp sin i ρp Planet Name a b −3 e ref. [days] [RJup] [MJup] [kg m ] [deg] OGLE-TR-56 b 1.212 1.30 ± 0.05 1.29 ± 0.12 720 0.00 81.0 22, 32, 38 TrES-3 1.306 1.295 ± 0.081 1.92 ± 0.23 1090 - 82.2 29 WASP-4 b 1.338 1.45 ± 0.06 1.27 ± 0.09 517 0.00 87.5 40 OGLE-TR-113 b 1.432 1.09 ± 0.03 1.32 ± 0.19 1240 0.00 89.4 7, 15 WASP-5 b 1.628 1.09 ± 0.08 1.58 ± 0.11 1519 0.00 > 85.0 2 OGLE-TR-132 b 1.690 1.18 ± 0.07 1.14 ± 0.12 850 0.00 85.0 7, 16 WASP-3 b 1.847 1.31 ± 0.14 1.76 ± 0.14 976 - 84.4 30 WASP-2 b 2.152 1.038 ± 0.050 0.88 ± 0.11 960 - 84.7 12, 13 HD 189733 b 2.219 1.154 ± 0.017 1.15 ± 0.04 910 0.00 85.7 8, 33 TrES-2 2.471 1.220 ± 0.045 1.198 ± 0.053 810 0.00 83.9 28, 37 WASP-1 b 2.520 1.443 ± 0.039 0.87 ± 0.07 350 - > 86.1 12, 13, 36 XO-2 b 2.616 0.97 ± 0.03 0.57 ± 0.06 750 - 88.6 9 GJ 436 b 2.644 0.37 ± 0.02 0.071 ± 0.006 1710 0.16 85.9 17, 18 HAT-P-5 b 2.788 1.26 ± 0.05 1.06 ± 0.11 650 0.00 86.8 5 HD 149026 b 2.877 0.73 ± 0.03 0.330 ± 0.023 1060 0.00 85.8 11, 35 HAT-P-3 b 2.899 0.89 ± 0.05 0.60 ± 0.03 1040 0.00 87.2 39 TrES-1 3.030 1.081 ± 0.029 0.75 ± 0.07 720 0.14 88.4 1, 43 HAT-P-4 b 3.056 1.27 ± 0.05 0.68 ± 0.04 400 0.00 89.9 24 OGLE-TR-10 b 3.101 1.24 ± 0.09 0.61 ± 0.13 390 0.00 89.2 23, 32 XO-3 b 3.192 1.92 ± 0.16 13.24 ± 0.64 2363 0.22 79.1 20 HD 209458 b 3.525 1.320 ± 0.025 0.657 ± 0.006 350 0.015 86.9 10, 21, 41 TrES-4 3.554 1.674 ± 0.094 0.84 ± 0.10 220 - 82.8 25 OGLE-TR-211 b 3.677 1.36 ± 0.18 1.03 ± 0.29 512 0.00 > 82.7 46 HAT-P-6 b 3.853 1.330 ± 0.061 1.057 ± 0.119 559 0.00 85.5 27 XO-1 b 3.942 1.18 ± 0.04 0.90 ± 0.07 660 - 89.3 19, 26 OGLE-TR-182 b 3.979 1.13 ± 0.24 1.01 ± 0.15 871 0.00 85.7 34 OGLE-TR-111 b 4.014 1.067 ± 0.054 0.53 ± 0.11 530 0.00 88.1 31, 42 HAT-P-1 b 4.465 1.203 ± 0.051 0.53 ± 0.04 370 0.00 86.2 3, 45 HAT-P-2 b 5.633 0.98 ± 0.04 8.04 ± 0.40 10400 0.50 > 86.8 4, 44 HD 17156 b 21.20 1.15 ± 0.11 3.12 2500 0.67 87.9 6, 14

1 Alonso et al. (2004), 2 Anderson et al. (2008), 3 Bakos et al. (2007a), 4 Bakos et al. (2007b), 5 Bakos et al. (2007c), 6 Barbieri et al. (2007), 7 Bouchy et al. (2004), 8 Bouchy et al. (2005), 9 Burke et al. (2007), 10 Charbonneau et al. (2000), 11 Charbonneau et al. (2006), 12 Charbonneau et al. (2007), 13 Collier Cameron et al. (2007a), 14 Fischer et al. (2007), 15 Gillon et al. (2006), 16 Gillon et al. (2007a), 17 Gillon et al. (2007b), 18 Gillon et al. (2007c), 19 Holman et al. (2006), 20 Johns-Krull et al. (2008), 21 Knutson et al. (2007a), 22 Konacki et al. (2003), 23 Konacki et al. (2005), 24 Kov´acs et al. (2007), 25 Mandushev et al. (2007), 26 McCullough et al. (2006), 27 Noyes et al. (2008), 28 O’Donovan et al. (2006), 29 O’Donovan et al. (2007), 30 Pollacco et al. (2007), 31 Pont et al. (2004), 32 Pont et al. (2007a), 33 Pont et al. (2007b), 34 Pont et al. (2007c), 35 Sato et al. (2005), 36 Shporer et al. (2007), 37 Sozzetti et al. (2007), 38 Torres et al. (2004), 39 Torres et al. (2007), 40 Wilson et al. (2008), 41 Winn et al. (2005), 42 Winn et al. (2007a), 43 Winn et al. (2007b), 44 Winn et al. (2007c), 45 Winn et al. (2007d), 46 Udalski et al. (2007) a 7 1RJup = 7.1492 ×10 m (equatorial radius) b 27 1MJup = 1.8986 ×10 kg

13 Chapter 2. Motivation, context and objectives

Table 2.2: Radius, mass and density of the giant planets of our solar system.

Rp Mp ρp Planet Name −3 [RJup] [MJup] [kg m ] Jupiter 1.000 1.000 1326 Saturn 0.843 0.299 687 Uranus 0.358 0.046 1270 Neptune 0.346 0.054 1638

from the planets of our solar system. In one hand, the close proximity of these objects to their stars has important consequences for their radius given the effects of tidal interactions and the intense stellar irradiation. HD 209458 b has a mean density that is almost a half that of Saturn, and a radius 10-20% larger than the expected for a planet with its age, mass, and temperature, even if we include effects of stellar irradiation (the stellar flux incident upon HD 209458 b is about 20 000 times that upon the surface of Jupiter), indicating the need for an additional source of energy to slow the evolutionary contraction of the planet. The models computed by Guillot et al. (1996) could explain the large radius of HD 209458 b, but only assuming unrealistic hot temperatures at deep layers. I has also been suggested (Guillot & Showman 2002; Baraffe et al. 2003) that, if a fraction of 1% of the stellar flux irradiating the planet is transported to deep layers, the internal heating produced could keep the radius as large as it is expected (see Figure 2.2). In contradiction to this scenario, the transiting planet TrES-1 (Alonso et al. 2004) have similar mass and similar distance to its parent star when compared to HD 209458 b, but its radius is not inflated. Clearly, the discussion is still open, and needs further investigation.

On the other hand, the transiting hot Saturn orbiting HD 149026 has a radius smaller than the predicted by standard models of giant planet formation (Bodenheimer et al. 2003). Such a small radius indicates that this planet is made of a significant fraction of heavy elements, about 50-90 M⊕ (see e.g. Fortney et al. 2006; Ikoma et al. 2006). For comparison, Ikoma et al. (2006) determined for HD 149026 b the same amount of hydrogen and helium as that of Saturn, but at least twice the mass of heavy elements. Saturn has a slightly smaller mass, is about 2.5 Gyr older, is irradiated with a stellar flux 140,000 smaller and, in spite of this, it has a mean radius larger than that of HD 149026 b.

Burrows et al. (2007) made a discussion concerning the observed radius of the known transiting extrasolar giant planets. They calculated theoretical evolution of their radii with the aim of understanding why some of these planets have radius larger than expected, whereas others are smaller than predicted by standard models. Their models propose that, in one hand, enhanced atmospheric opacities could explain the large radius, retarding the loss of heat and delaying the radius contraction. On the other hand, the presence of ice and/or rock cores could account for the observed small radius. The authors noticed, however, that their conclusions are based on several conditions.

14 2.3. Planetary transits and the search for hot Jupiters

Table 2.3: Parameters of the 30 stars with known transiting planets.

Spec. R! M! Teff Age dist. Star Name V a b [Fe/H] ref. type [R"] [M"] [K] [Gyr] [pc]

HD 209458 7.65 G0 V 1.13 ± 0.02 1.10 ± 0.07 6117 0.02 > 2 47 14, 17, 22 HD 189733 7.67 K1-2 0.75 ± 0.03 0.8 ± 0.4 5050 −0.03 > 0.5 19 6, 22, 36 HD 149026 8.15 G0 IV 1.5 ± 0.1 1.3 ± 0.1 6147 0.36 2 79 30 HD 17156 8.17 G5 1.47 ± 0.13 1.2 ± 0.1 6079 0.24 5.7 78 10 HAT-P-2 8.71 F8 1.48 ± 0.05 1.3 ± 0.1 6290 0.12 2.6 135 4, 37 XO-3 9.80 F5 V 2.13 ± 0.21 1.41 ± 0.08 6429 −0.18 2.7 260 16 HAT-P-1 10.0 G0 V 1.2 ± 0.1 1.12 ± 0.09 5975 0.13 3.7 139 3 HAT-P-6 10.5 F 1.46 ± 0.06 1.29 ± 0.06 6570 −0.13 2.3 260 23 WASP-3 10.5 F7-8 V 1.31 ± 0.09 1.24 ± 0.09 6400 0.00 - 220 26 GJ 436 10.66 M2.5 V 0.46 ± 0.02 0.44 ± 0.04 3350 0.00 > 3 10.1 13, 20 HAT-P-4 11.2 F 1.59 ± 0.07 1.26 ± 0.10 5860 0.24 4.2 310 18 XO-1 11.2 G1 V 0.93 ± 0.02 1.00 ± 0.03 5750 0.015 4.5 200 15, 21 XO-2 11.2 K0 V 0.97 ± 0.02 0.98 ± 0.02 5340 0.45 5.3 150 7 TrES-2 11.41 G0 V 1.00 ± 0.04 0.98 ± 0.06 5850 −0.15 5.1 220 24, 32 HAT-P-3 11.56 K 0.82 ± 0.04 0.94 ± 0.05 5185 0.27 0.4 140 33 TrES-4 11.59 - 1.74 ± 0.09 1.22 ± 0.17 6100 - 4.7 440 19 TrES-1 11.79 K0 V 0.83 ± 0.05 0.89 ± 0.05 5250 0.00 2.5 150 1, 31 WASP-1 11.79 F7 V 1.38 ± 0.08 1.24 ± 0.15 6200 - - - 8, 9 WASP-2 11.98 K1 V 0.83 ± 0.08 0.84 ± 0.09 5200 - - - 8, 9 HAT-P-5 12.0 G 1.17 ± 0.05 1.16 ± 0.06 5960 0.24 2.6 340 5 WASP-5 12.3 G4 V 0.97 ± 0.06 0.99 ± 0.08 5700 0.00 - 300 2 TrES-3 12.4 - 0.80 ± 0.05 0.90 ± 0.15 5720 - - - 25 WASP-4 12.5 G7 V 1.15 ± 0.28 0.90 ± 0.07 5500 0.00 - 300 35 OGLE-TR-211 14.3 c - 1.64 ± 0.14 1.33 ± 0.05 6325 0.11 - - 34 OGLE-TR-113 14.4 c K 0.77 ± 0.02 0.78 ± 0.02 4804 0.15 > 0.7 553 11, 22, 29 OGLE-TR-10 14.9 c G, K 1.14 ± 0.09 1.10 ± 0.05 6075 0.28 > 1.1 1326 22, 27, 29 OGLE-TR-56 15.3 c G 1.32 ± 0.06 1.17 ± 0.04 6119 0.25 > 2 1591 22, 27, 29 OGLE-TR-111 15.5 c G, K 0.83 ± 0.03 0.81 ± 0.06 5044 0.19 > 1.1 1011 22, 29, 38 OGLE-TR-132 15.7 c F 1.34 ± 0.08 1.26 ± 0.03 6210 0.37 - 2180 12, 29 OGLE-TR-182 15.9 c - 1.14 ± 0.15 1.14 ± 0.05 5924 0.37 - - 28

1 Alonso et al. (2004), 2 Anderson et al. (2008), 3 Bakos et al. (2007a), 4 Bakos et al. (2007b), 5 Bakos et al. (2007c), 6 Bouchy et al. (2005), 7 Burke et al. (2007), 8 Collier Cameron et al. (2007a), 9 Collier Cameron et al. (2007b), 10 Fischer et al. (2007), 11 Gillon et al. (2006), 12 Gillon et al. (2007a), 13 Gillon et al. (2007c), 14 Henry et al. (2000), 15 Holman et al. (2006), 16 Johns-Krull et al. (2008), 17 Knutson et al. (2007a), 18 Kov´acs et al. (2007), 19 Mandushev et al. (2007), 20 Maness et al. (2007), 21 McCullough et al. (2006), 22 Melo et al. (2006), 23 Noyes et al. (2008), 24 O’Donovan et al. (2006), 25 O’Donovan et al. (2007), 26 Pollacco et al. (2007), 27 Pont et al. (2007a), 28 Pont et al. (2007c), 29 Santos et al. (2006), 30 Sato et al. (2005), 31 Sozzetti et al. (2004), 32 Sozzetti et al. (2007), 33 Torres et al. (2007), 34 Udalski et al. (2007), 35 Wilson et al. (2008), 36 Winn et al. (2007a), 37 Winn et al. (2007c), 38 Winn et al. (2007e) a 1R" = 695,500 km (equatorial radius) b 30 1M" =1.9891 × 10 kg c apparent magnitude in the I band

15 Chapter 2. Motivation, context and objectives

HAT-P-2 b (Bakos et al. 2007b) and GJ 436 b (Gillon et al. 2007b) have very special characteristics in comparison to the other known transiting exoplanets. HAT-P-2 b has −3 large eccentricity (0.5), large mass (8.2 MJup), and large mean density (ρp = 6600 kg m ), which put it in an intermediate position between Jupiter-like planets and brown dwarfs (see Figure 4.2 in Section 4.2). GJ 436 b is the closest (10.2 pc), smallest (3.95 0.35 R⊕), ± and least massive (22.6 1.9M⊕) transiting planet detected so far. Measurements of ± the primary transit of GJ 436 b, gathered by Gillon et al. (2007b) using the Euler 1.2-m telescope at La Silla, and by Gillon et al. (2007c) using the 8 µm band of the InfraRed Array Camera (IRAC) on the Spitzer Space Telescope (Wernet et al. 2004), revealed that the planetary radius (0.37 RJup) is comparable to that of Neptune and Uranus, and suggests that the planet is mainly composed of water ice, with an H/He envelope.

Most recently, another transiting planet has been unveiled orbiting the star HD 17156 (Barbieri et al. 2007), with large eccentricity (0.67) and large orbital period (21.2 days). This planet has a favourable argument of periastron that combined with the large eccentric- ity provide a high transit probability. HD 17156 b was first discovered by the radial-velocity technique (Fischer et al. 2007), and was also confirmed by our radial-velocity measurements (see Figure 4.3). Such a large value for the eccentricity leads to a large variation of inso- lation between periastron and apoastron, which is an opportunity to study the dynamical flows in the planetary atmosphere. Probably, the planet has been experimenting significant tidal heating, but the derived radius is not inflated as one would expect.

Besides an accurate determination of the planetary radius and mean density, more precise estimates of other physical parameters of the system such as spin-orbit alignment, temperature at the planet’s surface, molecular and atomic composition of its atmosphere, cloud properties, albedos, among others, will provide important constraints on models of formation, interior structure, atmosphere, and orbital evolution of extrasolar planets. Bright stars offer an unique opportunity to perform this task by doing follow-up obser- vations with a large variety of techniques. In the next subsections I describe the main techniques and instruments currently used, and the published results for some of the known transiting systems (HD 189733 b is one of the results of this work, and a detailed discussion is done in Sections 4.1 and 4.2).

2.3.2 Transmission spectroscopy

This technique compares the observations gathered during the in-transit to those made out-of-transit to search for features in the transmission starlight through the outer parts of the planet’s atmosphere. As the stellar flux passes through the atmosphere of the planet, additional absorption features, associated with the presence of certain atoms or molecules in the planetary atmosphere, are superimposed to the spectral lines of the star. Hot Jupiters are expected to have extended atmospheres as a consequence of their high temperature, low gravity, and composition (mainly light gas), resulting in a cross section significantly large.

16 2.3. Planetary transits and the search for hot Jupiters

Charbonneau et al. (2002), using the Space Telescope Imaging Spectrograph (STIS) of the Hubble Space Telescope (HST), conducted in- and out-of-transit observations of the HD 209458 system, and detected the presence of gaseous sodium in the atmosphere of HD 209458 b. They compared the transit depth in the region of the Na absorption doublet centred at 5893 A˚ to that in adjacent bands, and measured a significantly deeper transit (by about 2 10−4), concluding that the additional dimming is due to the absorption × from this element in the atmosphere of the planet. Vidal-Madjar et al. (2003), using HST observations of the ultraviolet Lyman α emission line at 1215.67 A,˚ were able to detect a decrease of 15 4% in the intensity of this line during the transit, which indicates the ± presence of atomic Hydrogen in the outer atmosphere of HD 209458 b. Vidal-Madjar et al. (2004) made observations with the HST in the wavelength range λλ1180 1710, which − includes several spectral lines like C I, C II, C IV, N V, O I, S I, SiII, SiIII, and SiIV. They confirmed the detection of H I made by Vidal-Madjar et al. (2003) and, in addition, could detect absorption of O I and C II, indicating that such elements are also present in the planet’s atmosphere.

Molecules of H2O, CO and CH4 are also expected to be very abundant species in the atmosphere of close-in extrasolar planets. Barman (2007), by modelling the spectral emission and absorption in the atmosphere of HD 209458 b, concluded that his model is consistent with the presence of strong water absorption near1µm. In particular, the CO molecule is a diagnostic of the temperature in the planet’s atmosphere, since carbon migrates from CO to CH4 when the temperature is about 1400 K. Brown et al. (2002), by observing a transit of HD 209458 b using the NIRSPEC instrument connected to the Keck II telescope in the wavelength range 2.0 2.5 µm, detected no extra absorption of the − CO band near 2.3 µm. However, the observing conditions were not good enough to sustain their results. Deming et al. (2005b) have revisited the search for carbon monoxide in the atmosphere of HD 209458 b. They observed the transmission spectrum during transits of this planet in the region of the K-band around 2 µm, also using the NIRSPEC instrument. Again, no absorption due to the presence of CO in the planet’s atmosphere was detected. They suggested that, as already pointed out by Richardson et al. (2003b), CO molecules presumed to be present in the atmosphere of HD 209458 b may effectively be hidden by cloud opacity.

The presence of signature of other elements in the transmission spectra of exoplan- ets have also been investigated. Bozorgnia et al. (2006) observed the transiting planet HD 149026 b using the Keck I Telescope before and during transit, and a differential com- parison of the spectra detected no signatures of neutral atomic lithium or potassium in its atmosphere.

2.3.3 Occultation photometry and spectroscopy

In this approach, observations are made during secondary eclipses, or anti-transits (when the planet passes behind the star) and then compared to those made just before

17 Chapter 2. Motivation, context and objectives

or just after the anti-transit. This provides an investigation of the reflected light from the planet, which is sensitive to the nature and composition of its atmosphere. Such observations are normally conducted in the near- and mid-infrared, where the planet-to- star flux ratio is most favourable. Burrows et al. (2006) computed cloud-free models for five close-in exoplanets, showing a variation of more than 4 orders of magnitude from the optical to the mid-infrared. Therefore, this technique offers the possibility to search for any excess in thermal radiation emitted by the planet.

In our solar system, water vapour can exist only in planetary atmospheres at orbital distances smaller than 1 AU. Water vapour is thus expected to be present in the atmosphere of close-in giant exoplanets. So are the carbon monoxide, carbon dioxide, methane, and others. A model computed by Tinetti et al. (2007) predicts absorption signatures due to water vapour in the atmospheres of HD 209458 b and HD 189733 b between 3 and 25 µm. Deming et al. (2005a), using the Spitzer telescope at 24 µm, conducted observations during a secondary eclipse of the HD 209458 system. They computed a model atmosphere for HD 209458 b that is compatible with the presence of strong water vapour absorption at this wavelength.

Richardson et al. (2003a,b) conducted infrared observations during secondary eclipses of the HD 209458 system at the methane band near 3.6 µm, using the ISAAC spectrometer on the Very Large Telescope (VLT), and at the CO and H2O absorption bands near 2.2 µm, using the SpeX instrument at the NASA Infrared Telescope Facility (IRTF). Their results, however, show no evidence for the presence of such infrared signatures despite the theoretical predictions, and put new constraints on models that describe the atmosphere of giant planets.

The fraction of the incident stellar radiation scattered by a planet depends on the parameter called albedo. During anti-transit events, the albedo is defined as the ratio of the planet’s luminosity at full phase to the luminosity from a Lambert disk2 with the same cross-section as the planet. It is an important constraint for theoretical model atmospheres, excluding or including the presence of reflective clouds. Rowe et al. (2006), using observations from the Microvariability and Oscillations of Stars (MOST) satellite (Walker et al. 2003; Matthews et al. 2004), determined an upper limit on the albedo of HD 209458 b (" 0.25). This is a small value compared to the planets of the solar system, which have an albedo ! 0.4 in the MOST bandpass (400-700 nm). The solar system giant planets all have bright clouds of water and ammonia ice, but HD 209458 b is too hot to allow such a composition.

Another quantity that is associated to anti-transit observations is the orbital eccen- tricity of the planet. Effects of non-zero can change the time of the centre of secondary eclipses relative to primary eclipses, producing a shift that deviates the difference between them away from a half-period (see Charbonneau 2003).

2 A Lambertian surface is an ideal, isotropic reflector at all wavelengths.

18 2.3. Planetary transits and the search for hot Jupiters

Figure 2.3: Examples showing the dependence of the Rossiter-McLaughlin waveform on the angle λ , for λ = 0, 30, and 60 ◦. The impact parameter b is fixed at 0.5. The so so − dotted lines represent the case of no limb darkening, whereas the solid lines are for the case of a linear law. Figure from Gaudi & Winn (2007).

2.3.4 The Rossiter-McLaughlin effect

When the planet (or the stellar companion in the case of a binary system) crosses the disk of the main star during a transit event, it hides some of the rotational velocity fields that contributes to the line broadening. The result is a distortion of the spectral line profiles, which leads to an additional shift in the radial velocities that is superimposed to the radial-velocity variation induced by the orbital movement of the companion. It was first detected on the eclipsing binaries β Lyrae and Algol by Rossiter (1924) and McLaughlin (1924), and first observed in an exoplanetary system by Queloz et al. (2000) and Bundy & Marcy (2000) during transits of HD 209458 b. The Rossiter-McLaughlin effect depends on the projected rotation velocity of the stellar surface, and on the angle λso between the sky projections of the stellar spin axis and the orbit normal. Figure 2.3 shows some examples of what the distortion in the Keplerian curve looks like, due to the additional radial-velocity shift, for different values of λso. The consequences of including the limb-darkening effect is also represented in this figure.

For the planets of the solar system, the angle between their orbital plane and the solar equatorial plane is generally within 5 ◦. In the case of the close-in giant planets ∼ in the extrasolar planetary systems, they are believed to have formed at large orbital distances and then migrated to their present orbit, closer to their host star. Thus, one might wonder if the migration process would be able to disturb the alignment of the spin-orbit axes. By studying and modelling the Rossiter-McLaughlin effect, one have the possibility of measuring such an angle, and verifying if the spin-orbit alignment is also

19 Chapter 2. Motivation, context and objectives

valid for extrasolar planetary systems (see Gaudi & Winn 2007).

So far, many authors have determined the angle λso for some of the known transiting systems by doing observations of the Rossiter-McLaughlin effect, and they all found small values, consistent with zero. For the HD 209458 system, Queloz et al. (2000) and Winn ◦ ◦ et al. (2005) found λso =3.9 21 and λso =4.1 1.4 , respectively. Wolf et al. (2007) ± ± ◦ observed the HD 149026 system, and derived a value of λ = 12 15 . Winn et al. so − ± (2007c) and Loeillet et al. (2007) derived small values for such an angle in the HAT-P-2 system, which are also consistent with zero within about 14◦ and 9 ◦, respectively. Narita et al. (2007) measured λ = 30 21 ◦ for the TrES-1 transiting system, which is larger than so ± the measurements obtained for the other systems, but is also consistent with zero taking into account the large error. Given such results, it seems that a spin-orbit misalignment is not achieved at the end of the inward migration process, suggesting that the planet-disk interaction does not significantly perturb the coplanarization. Nevertheless, one need more measurements of this angle, with better precision, and for a larger number of systems, so that statistical studies can be performed.

2.3.5 The transit light curve

The observed light curve from a star with a transiting planetary companion can be generally described, under simplifying assumptions, by a set of five equations. Three of these equations describe the geometry of the transit event: depth, shape, and duration (Equations 2.2 to 2.4). The transit depth ∆F is related to the planetary radius Rp and to the stellar radius R!. The transit shape is a relation between the total transit duration tT and the flat transit duration tF between ingress and egress. The total transit duration is the time from the first to the fourth contact between the planet and the stellar disk, while the flat transit duration is the time when the planet is fully in front of the stellar disk, between second and third contacts. See Figure 2.4 for an illustration of the geometrical and physical quantities in question. F F R 2 ∆F = out of transit − in transit = p (2.2) F R out of transit ! ! "

sin (t π/P ) [1 (R /R )]2 [(a/R ) cos i]2 1/2 F = p ! ! (2.3) { − 2 − 2}1/2 sin (tTπ/P ) [1 + (R /R )] [(a/R ) cos i] { p ! − ! }

P R [1 + (R /R )]2 [(a/R ) cos i]2 1/2 t = arcsin ! p ! − ! (2.4) T π a 1 cos2 i % & − ' (

In the equations above (which are from Seager & Mall´en-Ornelas 2003), P , a and i are the canonical orbital parameters: period, semi-major axis and inclination angle, respec- tively. The observables that come directly from the transit light curve are ∆F , tT and tF.

20 2.3. Planetary transits and the search for hot Jupiters

R*

Rp132 4

bR* = a cosi

1 2 3 4

!F

t F t T

Figure 2.4: Schematic illustration of a simplified transiting light curve. Solid and dotted lines represent the light curve shapes for two different values of impact parameter. Posi- tions 1 and 4 indicate the first and fourth contacts, respectively, and the difference in time between them gives the total transit duration tT. Positions 3 and 4 indicate the ingress and egress, respectively, and their difference in time is the flat transit duration tF. The transit depth ∆F is also indicated. Figure from Seager & Mall´en-Ornelas (2003).

From Equation 2.2 we derive the planet-to-star radius ratio Rp/R!; from Equation 2.3 we derive the impact parameter, defined as b =(a/R!) cos i; and from Equation 2.4 we derive the ratio a/R!. The other two equations required in this approximation for a complete, although simplified, description of the transit light curve are the Kepler’s third law for a circular orbit, and the stellar mass-radius relation:

4π2a3 P 2 = (2.5) G(M! + Mp)

x R! = kM! (2.6) where M! and Mp are, respectively, the stellar and planetary mass, G is the universal gravitational constant, and the coefficients k and x are constants that depends on the stellar position in the HR diagram (many authors adopt k = 1 and x 0.8 for main- & sequence stars).

If we assume M M , we can rewrite Equation 2.5 as the ratio a3/M = GP 2/4π2. p ' ! ! We have then four relations that depend only on the observables, and together with the stellar mass-radius relation they provide the determination of the five physical parameters describing the transit light curve: R!, M!, Rp, i, and a.

21 Chapter 2. Motivation, context and objectives

The description above of the transit light curve was made under several assumptions aiming to present, in a simplified way, the relationship between observed and derived geo- metrical and physical parameters in question. In summary, the simplifying considerations are the following:

i) the orbit is circular: this is a good approximation since short-period planets are thought to have their orbits circularised by tidal interactions (Halbwachs et al. 2005). Indeed, more than 80% of the discovered extrasolar planets with periods shorter than 5 days have e<0.1, taken into account only those with known eccentricity.

ii) the planetary radius is much smaller than the stellar radius: such an approximation should be used carefully. Equation 2.2 shows that larger is the planet-to-star radius ratio deeper is the dimming in the transit light curve, and thus the transit event is easier detected. Most of the known transiting systems have this ratio around the limit of 10%, or even larger in some cases.

iii) the adoption of a stellar mass-radius relation: although the limitations to derive the mass of the star based on models of stellar evolution impose uncertainties in the determination of the radius of the planet, this approach is required if the quality of the light curve is not good enough to assure an independent estimate of this parameter.

iv) the orbital period of the planet is known: the period can be determined from the light curve if two consecutive transits are observed. Otherwise, if radial-velocity measurements are available, the period can be obtained from the Keplerian solution fitted to the velocity data.

v) the intensity of the radiation emitted by the star is constant throughout the stellar disk: in practice, this approximation may be done when dealing with low precision data, or just to have a first estimate of the transit parameters. However, a more real- istic treatment of the situation should take the so-called limb-darkening phenomenon into account (see below). In particular, for wavelengths longer than about 1 µm the effect of the limb darkening is weak, and can be neglected (Knutson et al. 2007a). Otherwise, it should be considered if we want to have a more accurate estimate of the parameters, especially when precise photometric data are available.

Another important feature of the transit light curve fit concerns the degeneracy in Rp, R!, and i. If the planetary and stellar radius reduce (increase) in proportion, the transit depth is preserved. If, in addition, the orbital inclination reduce (increase) accordingly, the chord length across the stellar disk is conserved as well. Therefore, we can fit different models to the same observed light curve just by doing different combinations of these parameters. There are some situations for which it is possible to break this degeneracy. If high quality photometry is available so that the shape of the ingress and egress can be fitted with accuracy, then we can estimate i and Rp unambiguously. Furthermore, if we observe the transit light curve in multiple bandpass, then the inclination can be uniquely

22 2.3. Planetary transits and the search for hot Jupiters

(a) (b) observer rr ** rr ** rr =pr pp ** "" d=zr ** rr rr planet ** µµ ** star star

Figure 2.5: The limb darkening (a) and the transit (b) geometry. The stellar and planetary radius are r! and rp, respectively. The z parameter is the normalised separation of the centres, p is the radius ratio, and µ = cos θ. Figure from Mandel & Agol (2002). determined (see e.g. Jha et al. 2000; Knutson et al. 2007a). In fact, combining the orbital information from the radial velocities with a very precise light curve can reduce the number of free parameters to only one: the stellar mass.

An analytical solution able to fit a given transit light curve is not unique if contami- nation by stellar blends is present. The light from the stars in a eclipsing binary system, combined to the light from a third star (physically bound or just a background star), can mimic the observed transit light curve of a planet (see discussion in Seager & Mall´en- Ornelas 2003). In such situations, radial-velocity measurements of the system together with a line bisector analysis (see Subsection 2.4.2) could unveil the stellar nature of the companions.

Limb-darkening effect

An important effect that should normally not be neglected when modelling a transit light curve is the limb darkening, which accounts for the fact that the intensity throughout the stellar disk is not uniform, decreasing from the centre to the edge (see Mandel & Agol 2002). For main-sequence stars, it is described by functions of the parameter µ = cos θ, where θ is the angle between the normal to the stellar surface and the line of sight (see Figure 2.5). Common laws proposed to describe the specific intensity throughout the stellar surface have linear, quadratic, square root, or even logarithmic dependence on µ. Claret (2000) proposed a new non-linear law, for which the specific intensity I(r) in each point of the stellar disk is given by:

4 I(r) = 1 c (1 µn/2) (2.7) − n − )n=1 where µ = cos θ = (1 r2)1/2,0 r 1, and c are coefficients to be determined. − ≤ ≤ n

23 Chapter 2. Motivation, context and objectives

In one hand, the nonlinear approach provides an accurate and more realistic charac- terisation of the limb darkening. On the other hand, it requires a high accuracy in the observations of the transit light curve so that the adjusted coefficients could be well con- strained. Another usual solution is the quadratic limb-darkening law, which still provides a relatively good accuracy. In this approach, the specific intensity is given by:

I(r) = 1 γ (1 µ) γ (1 µ)2 (2.8) − 1 − − 2 − where γ1 and γ2 are coefficients to be determined, and γ1 + γ2 < 1. The nonlinear law (Equation 2.7) reduces to the quadratic law when c = c = 0, c = γ +2γ , and c = γ . 1 3 2 1 2 4 − 2 In some situations it is possible to make additional approximations that still provide reasonable accuracy. This is the case when the planet is much smaller than the star (radius ratio p " 0.1), and we can then assume that the specific intensity is constant under the portion of the stellar disk occulted by the planet. The choice of one of the three approaches (nonlinear, quadratic or small planet) will depend on each specific light curve. We can have a fast first estimate of the transit parameters using the small planet approximation and then, depending on the quality of the photometric data, use either the quadratic or the nonlinear approach to refine the derived values.

2.4 Other sources of RV variability

The radial-velocity technique employed to the search for extrasolar planets is based on the detection of variations in radial velocity caused by the gravitational interaction between the observed star and the presumed orbiting planet. However, the origin of the observed radial-velocity variations can be related to other astrophysical phenomena rather than the orbital motion induced by the presence of a planetary companion. We can classify these sources of spurious radial-velocity in the following categories: 1) variability intrinsic to the star, such as those induced by the stellar pulsation, convective granulation, or chromospheric magnetic activity; 2) variability due to the proper motion of the star (secular acceleration); 3) variability induced by a faint stellar companion in a blended spectra; and/or 4) variability related to the Earth movement around the Sun.

Concerning the stellar pulsation, the modes of oscillation commonly observed in solar- type stars are p-modes, in which the restoring force is provided by the pressure. The variation induced in the radial velocities is typically of a few tens of cm s−1 over a few minutes, which is very small compared to the time scale of typical orbital periods (> 1 day). It is important in high-precision radial-velocity surveys, as is the case of planet search pro- grammes using the HARPS (High-Accuracy Radial-velocity Planet Search) spectrograph, mounted on the 3.6-m telescope at La Silla (ESO, Chile), for which we can attain a preci- sion smaller than 1 m s−1. Oscillations related to the convective granulation is also small, and is only important in high-precision studies. Another kind of oscillation is non-radial pulsation, associated with g-modes, in which the restoring force is buoyancy (involving

24 2.4. Other sources of RV variability

both gravity and composition gradients). However, it is unlikely to be significant in solar- like stars (see Brown et al. 1998). Taking into account the purposes of the present work, and the limitations of the ELODIE spectrograph, these sources of radial-velocity variability are dominated by other sources of uncertainties.

The secular acceleration is a change in the radial velocity of the stars due to their movement in the space. It depends on the proper motion, and on the distance of the star, being important for high proper motion stars at high distances. For the star from our sample, the secular acceleration is smaller than 1 m s−1/yr, and can be neglected.

The stellar spectrum as observed by ground-base telescopes is contaminated by telluric absorption lines (absorption features formed in the Earth’s atmosphere), especially in the red portion of the visible wavelengths. The relative position between the absorption lines in the stellar spectrum and the telluric lines depends on the relative movement between star and observer. This contamination introduce uncertainties in the computed CCF that can result in periodic changes in the radial velocities, especially an one-year period of variation related to the Earth’s movement around the Sun. To correct from this contamination, or at least minimise its effects, numerical masks have been constructed using only the orders with no detected telluric lines. An example presented in Naef (2004) concerns the star HD 145729, for which the one-year periodic signal, found using radial velocities determined with previous masks, disappears when using the ”cleaned” ones. An improvement in the precision of the radial velocities is also achieved with these masks. Naef (2004) lists some other examples of known systems with a substantial reduction of at least 30% in the rms around the adopted Keplerian solution. It should be noticed, however, that the effect of the telluric lines are not perfectly corrected, and residual periodic variations of 1 year may still exist. Moreover, remembering that we measure the radial velocities of the stars setting the solar system barycentre as reference, other residual variations with the same period may also exist as a result of an imperfect correction of the Earth’s movement.

In this work, the variability induced by a faint stellar companion, and that caused by the magnetic activity on the chromosphere of the star are the dominant sources of spurious radial-velocity variation, and are better discussed below. Both these phenomena can lead to significant asymmetries on the profiles of the spectral lines that will affect the radial-velocity measurements and induce erroneous interpretation.

2.4.1 The stellar chromospheric activity

The magnetic activity on the chromosphere of late-type stars is predicted by the dy- namo theory, where magnetic fields regenerated in the convection zone are ruled by the differential stellar rotation. In this way, higher rotation rate are expected to cause stronger magnetic activity, which in turn enhances the number of starspots present in the photo- sphere as well as the flow of material in active regions. Inhomogeneities in the starspots distribution throughout the stellar disk, and spatial variation of granulation in the convec-

25 Chapter 2. Motivation, context and objectives

Figure 2.6: Examples comparing the strong emission in the core of the CaII absorption line at λ3968.5 of the active star HD 166435 (left) to the same absorption line of the non- active star HD 68255 A (right), which is one of our programme stars, both of the same spectral type (G0 V). The value for log R& is 4.27 for the active star. HK − tive zone often affect the line shape in the spectrum of active stars, since they rotate across the disk and evolve with time. In such cases, the centroid of the line profiles are shifted, resulting in an apparently variable radial velocity. The measurement of the chromospheric activity level of a star is thus a useful tool to identify possible sources of contamination in the radial-velocity observations.

Several strong absorption lines in the stellar spectrum have been suggested to be used as a diagnostic of the magnetic activity level. In active stars, those lines exhibit an enhanced emission3 in the core due to chromospheric radiation contribution. The ultraviolet doublet of MgII h & k at λ2802.7 and λ2795.5, the infrared triplet of Ca II at λ8498.1, λ8542.1 and λ8662.2, and the Hα line of hydrogen at λ6562.8 are typical examples of such an application. The doublet of Ca II H & K at λ3968.5 and λ3933.7 is nevertheless the most recognised as a good indicator of the degree of magnetic activity on the chromosphere of late-type stars (see e.g. Noyes et al. 1984; Saar et al. 1998; Strassmeier et al. 2000). Figure 2.6 shows two spectra in the region of the Ca II H absorption line, comparing the core emission of an active star (HD 166435) to the absence of emission of a non-active star (HD 68255 A, one of our target stars).

A quantitative measurement of the chromospheric activity level using the Ca II H & K lines was performed by Vaughan et al. (1978). They introduced the Mount Wilson S index

3 The cause of this emission is the higher temperature of the chromosphere, where the core of strong lines is formed, in comparison to the photosphere, where their wings are formed.

26 2.4. Other sources of RV variability

of activity, defined as being proportional to the ratio of the measured flux within the 1 AH˚ and K bandpass of the Mount Wilson spectrophotometer to that in two continuum windows equidistant to the red and violet of the H and K lines. The S index is however dependent on spectral type, being affected by residual contributions of photospheric radiation at the centres of the H and K lines. When dealing with stars of different spectral types, the & parameter RHK is more often employed instead. It is defined as:

R& = R R (2.9) HK HK − phot where RHK is the ratio of the emission in the cores of the Ca II H and K lines to the total 4 bolometric emission of the star (FHK/σTeff ), and Rphot is the photospheric contribution of the stellar radiation to the observed H and K emission that is transmitted by the H and K passbands.

& The RHK index represents the mean level of activity, and are normally given in logarith- mic scale. Active stars have typical values of logR& between 4.75 and 4.2, the latter HK − − representing a higher activity level than the former. But there are some extreme cases of very active stars, like HD 129333 and HD 143006, with logR& = 4.18 and log R& HK − HK = 4.03, respectively (Wright et al. 2004). The Sun is slightly active, with a mean level − & log RHK = 4.94 (Hall et al. 2007). Typical quiet stars concerning their level of activity ( ) & − have log R smaller than 5.0. HK − It is important to notice that a strong level of chromospheric activity does not, by itself, imply that changes in the observed radial velocities of a star will occur. The positions of the active regions (where spots are located: pole, equator), the spatial distribution of a non-axisymmetric velocity field (related to convective inhomogeneities), and the rotation rate of the star (given by the v sin i value) are factors that affect the variability (see Desort et al. 2007). In the next section I present the method of the line bisectors, commonly used to distinguish magnetic activity effects from real radial-velocity variations due to the presence of a planetary companion.

2.4.2 Line bisector analysis

The observed profile of absorption lines in the solar and stellar spectrum has been largely studied in past years, aiming to have a better understanding of its behaviour and origin (see e.g. Asplund et al. 2000; Gray 2005, and references therein). One of their most investigated properties concerns the fact that each line profile shows an asymmetric structure relative to its central position in wavelength.

Asymmetries of spectral line profiles can originate from many different sources: os- cillation, pulsation, starspots on the surface, convective granulation. In cool stars, they originate specially from velocity fields on the granulated surface, in which material from different regions contributes with different shifts to the line formation process. Given the

27 Chapter 2. Motivation, context and objectives

Figure 2.7: Individual line bisectors of the Fe I at λ6252 for stars of different spectral types and luminosity classes. Crosses show the calculated points of the bisector with error bars on the velocity. Figure from Dall et al. (2006).

relation between convection and magnetic fields, such asymmetries have a direct depen- dence on the magnetic activity of stars, being strongly affected by spots rotating across the disk. The magnetic activity can thus induce periodical variations to the line asymme- tries. We can measure possible changes in the line asymmetries by computing the so-called line bisectors. The line bisector is constructed by finding and connecting the midpoint of horizontal line segments going from points on the left to interpolated points on the right side of the profile. Figure 2.7 shows the typical shape of line bisectors for stars of different spectral types and luminosity classes.

Originally, bisectors were calculated using single absorption lines from very high S/N and high resolution spectra. However, this normally requires the combination of several individual spectra obtained in different times, which may introduce errors. Dall et al. (2006) have investigated the behaviour of line bisectors derived from single spectral lines of combined spectra in comparison to those derived from CCF of single-exposure spectra. They confirmed the conclusion of Queloz et al. (2001) that line bisectors may be calculated from the CCF profile instead of individual lines, since the errors due to changes in time are eliminated when using one single spectrum.

Changes in line bisector orientations can be easily evaluated by computing the so-called

28 2.4. Other sources of RV variability

Figure 2.8: Left: The bisector inverse slope (BIS) of HD166435 versus its radial velocities. Right: BIS values and radial velocities phase-folded with the period P = 3.8 days. A clear anticorrelation is observed. bisector inverse slope (BIS), which is the difference between the mean velocity on the top of the bisectorv ¯t and that on the bottomv ¯b (Queloz et al. 2001). The main idea of this approach was already used by Mayor & Queloz (1995) in their analysis of the 51 Peg system, and it is similar to the bisector velocity span defined by Toner & Gray (1988). The velocitiesv ¯t andv ¯b are determined by averaging selected points on top and bottom, but neglecting points near the continuum and the core, since they may increase the noise.

In the case we observe a star moving as whole due to the perturbation of a companion, the computed line bisectors would oscillate around the mean bisector, with no change in shape or orientation, in such a way that the BIS value would remain constant. On the other hand, if the movement originates in the stellar surface, for example due to spots rotating around the star, the line bisectors would twist around the mean bisector, meaning that the BIS value would have a periodic variation correlated with the stellar rotation period. We can verify if there is any correlation between changes in line-bisector orientations and changes in radial velocities by plotting the BIS values against the radial velocities. A classical example is the case of HD 166435. The radial-velocity variations of this star were, at first, wrongly interpreted as the result of planetary reflex motion. After a series of photometric and Ca II H and K observations, the variations revealed to be induced by the surface magnetic activity of the star, in modulation with its rotation (Queloz et al. 2001). The left panel of Figure 2.8 shows the BIS values of HD 166435 plotted against the radial velocities. On the right panel of the same figure, radial velocities and BIS values are plotted phase-folded with the period determined by the Keplerian fit ( 3.8 days, which ∼ is equal to the period derived from the photometry). Clearly, we can observe that both variations are in perfect anticorrelation.

29 Chapter 2. Motivation, context and objectives

Figure 2.9: Left: Illustration of the effect of changing the relative position of HD 41004 A (deeper profile) and HD 41004 B (smaller profile) spectra. The shape of the blended profile is asymmetric as shown by the line bisector (dashed lines). Right: BIS values of HD 41004 versus its radial velocities. The plot shows that the two quantities are correlated. Figures from Santos et al. (2002).

The line bisector analysis has also been used to identify possible contamination by an unseen companion. An example concerns the system HD 41004, which is composed of four bodies: a primary K1 star orbited by a 2.3 MJup planet (Zucker et al. 2004), and a secondary M2 star with a brown-dwarf companion of 18.4 MJup (Zucker et al. 2003). The planet orbits the main star with a long period of 2.6 yr, while the brown dwarf has a very short-period orbit of only 1.3 days. HD 41004 A and B are separated by only 0.5&&, and we actually observe the blended spectra of the two stars. By doing a bisector analysis, complemented with photometric observations, Santos et al. (2002) proposed that the presence of a brown-dwarf orbiting the fainter companion is the best explanation for the observed short-term radial-velocity variation. In this scenario, the spectrum of the faint star is shifted in wavelength due the presence of the brown-dwarf, and oscillates relative to the spectrum of the main star. The resulted blended lines thus have an asymmetric profile, in which the top of the line bisector oscillates around the mean bisector (Figure 2.9, right panel). The effect is the opposite to that observed when the line-asymmetry variations are due to magnetic-related phenomena: the top of the line bisector oscillates around the mean bisector, and we observe a correlation between BIS and radial velocities (Figure 2.9, left panel) instead of an anticorrelation.

30 Chapter 3

Sample selection and observations

Contents

3.1 Selection of the target stars ...... 31 3.2 Observations ...... 33 3.3 Calibrations using the CCF profile ...... 35 3.4 Global view of the observations ...... 40

3.1 Selection of the target stars

We created our sample using solar-type stars chosen from the Hipparcos catalogue (ESA 1997) according to the following conditions:

i) stars in the northern hemisphere (declination δ 0), within a distance 100 pc ≥ ≤ (parallax π 10 mas), and visual magnitude V 8.5. ≥ ≤ ii) F8 spectral type < M1, corresponding to the colour index range 0.45 < (B V ) ≤ − ≤ 1.4. Stars hotter than the solar-type ones have smaller number of spectral lines and broader line profiles. These effects would cause uncertainties in the radial velocities, since they are determined by the CCF, which depends on the individual shape of the spectral lines. On the other hand, M stars are too faint to provide good precision in the radial velocities when observed with the ELODIE spectrograph. I searched for the (B V ) colour index in the SIMBAD Astronomical Database1 in the case of − some stars with no available value in the Hipparcos catalogue. A total of 3340 stars were selected at this point.

1 http://simbad.u-strasbg.fr/simbad/

31 Chapter 3. Sample selection and observations

iii) we eliminated the evolved stars, selecting only those within 1.5 mag above the main sequence2. Evolved stars are well known for their intrinsic radial-velocity variation due to pulsation and/or jitter. Stars located below the main sequence (probably due to errors in their parallaxes) should be excluded as well, but we found no similar case among those in our pre-selection. A total of 2605 stars selected up to here.

iv) comparing our pre-selected sample with the main surveys of extra-solar planetary search, we excluded common stars between our sample and those from previous pro- grammes that were based on observations collected with the ELODIE spectrograph or the Keck telescope. We found 543 stars already present in other programmes, remaining 2062 objects.

v) finally, using the Hipparcos data together with the CORAVEL database of radial and rotational velocities (Benz & Mayor 1984; Duquennoy et al. 1991; Duquennoy & Mayor 1991), we identified and excluded binary stars and fast-rotating stars. That includes: spectroscopic binaries of two spectra (the so-called SB2); close binaries with angular separation within 10 arcseconds; stars with errors in radial velocity greater than 2 km s−1; and stars with high projected rotation velocity (v sin i>6 km s−1). High-rotation stars have spectra with broad line profiles, producing errors in the CCF fit. Moreover, the rotational velocity is directly related to the chromospheric activity level of main-sequence stars (Noyes et al. 1984), in the sense that stars with higher rotation are more active chromospherically. The occurrence of physical events in the atmosphere of active stars, like the presence of spots, can induce spurious radial- velocity variations. Similar effects are expected in the case of spectra with blended lines generated by the presence of a unresolved companion in a multiple system. This step eliminated 1001 stars from those selected just before.

A catalogue containing a total of 1061 solar-type stars was then created (see Table A.1). Figure 3.1 shows the HR diagram (left panel) and the colour-distance diagram (right panel) for all the selected stars, using the (B V ) colour. − In an a posteriori initiative, we decided to create two subsamples of stars with small declination to be observed in the southern hemisphere. One of these subsamples consists of 20 stars with [Fe/H] 0.0 and declination between 0 and +20 ◦, which we selected to ≥ be monitored using the HARPS spectrograph as part of an observation programme whose principal investigator was C.H.F Melo from the Department of Astronomy of the Chile University. The main goal was to make use of the high precision achieved with HARPS to search for very low mass extrasolar planets with short periods. The other subsample consists of 98 stars with [Fe/H] 0.1 and declination between 0 and +10 ◦, selected to ≥− be observed using the CORALIE spectrograph. This subsample is still part of our main programme, but was created aiming to monitor low-declination stars after the end of the ELODIE spectrograph (see next section).

2 We used the main sequence as described by the relation MV = 1 + 5.4(B − V ), which represents an averaged main sequence defined by Naef (2004) based on a sample of stars observed with CORALIE.

32 3.2. Observations

Figure 3.1: HR diagram (left) and colour-distance diagram (right) for the 1061 selected stars. The dashed line on the left panel indicates the limit in magnitude chosen to eliminate evolved stars (see text).

3.2 Observations

We started the observations of our target stars using the high-precision ELODIE fibre-fed echelle spectrograph (Baranne et al. 1996), mounted on the 1.93-m telescope at the Haute-Provence Observatory (France). This instrument yields a spectral resolution (λ/∆λ) of 42 000. For typical exposure times of 20 min we had spectra with signal-to-noise ratio (S/N) between 20 and 100, which correspond to photon-noise errors between 5 − ∼ ∼ and 20 m s 1 on single measurements, mainly depending on the stellar magnitude. The observations were conducted from March 2004 to August 2006, in a total of 220 nights distributed over 29 observing missions (see Table 3.1 for more details). About 20% of these nights was lost due to bad weather conditions or technique problems. Moreover, the observing time allocated to our programme was shared with other ELODIE planet search programmes (about 16% of the 220 nights was dedicated to such programmes). Thus, about 64% of the allocated time for the observations was dedicated exclusively to our programme, which corresponds to approximately 140 nights. In the course of these nights, we collected a total of 2341 spectra of 854 observed stars.

The observational strategy consisted in getting at least one measurement of each star, estimating its metallicity and monitoring only those with [Fe/H] 0.1. Stars with v sin i − ≥ > 6 km s 1 were not taken into account as well, given the non-gaussian profiles of fast rotators. We used the ELODIE CCF to determine the stellar metallicity and projected rotation velocity, as described in the following section . After doing a minimum of three measurements of the metal-rich stars, those showing a significant variation in the radial

33 Chapter 3. Sample selection and observations

Table 3.1: Number of nights of the observing missions that were allocated to our main programme using the ELODIE spectrograph. From the total observing time (220 nights), about 20% was lost due to bad weather conditions or technique problems, and about 16% was dedicated to other ELODIE planet search programmes. Allocated Allocated Observing missions Observing missions nights nights 2004-03-01 2004-03-07 7 2005-08-27 2005-08-31 5 → → 2004-05-24 2004-05-31 8 2005-09-09 2005-09-15 7 → → 2004-07-26 2004-08-08 14 2005-09-27 2005-09-29 3 → → 2004-09-26 2004-10-03 8 2005-10-23 2005-10-25 8 → → 2004-10-25 2004-10-27 3 2005-11-11 2005-11-18 3 → → 2004-11-22 2004-11-29 8 2005-12-09 2005-12-23 15 → → 2005-01-21 2005-01-28 8 2006-01-12 2006-01-23 12 → → 2005-02-19 2005-02-28 10 2006-02-06 2006-02-14 9 → → 2005-03-30 2005-04-04 6 2006-03-12 2006-03-18 7 → → 2005-04-24 2005-05-02 9 2006-04-13 2006-04-20 8 → → 2005-05-25 2005-05-31 7 2006-15-13 2006-05-19 7 → → 2005-06-14 2005-06-19 6 2006-06-07 2006-06-14 8 → → 2005-06-27 2005-07-03 7 2006-07-15 2006-07-19 5 → → 2005-07-22 2005-07-25 4 2006-08-11 2006-08-13 3 → → 2005-08-01 2005-08-15 15 → Total of allocated nights: 220 Nights dedicated to our main programme: 140 ∼ velocities (! 50 m s−1) were monitored regularly to investigate the possibility of periodic variations. Since we aimed to discover planets with short-period orbits (shorter than 10 days), we tried to make the first two measurements of the same star separated by two days. Doing so, we expected to increase our chances to detect radial-velocity variations involving short-time scales.

As discussed in Section 2.3, planets orbiting their parent stars with short periods are good candidates to search for planetary transits. For this reason, a complementary programme of photometric observations was initiated using the 1.2-m telescope at the Haute-Provence Observatory. Hot Jupiters eventually discovered would be photometrically monitored during the predicted time for the transit event.

The ELODIE spectrograph was decommissioned in August 2006, before the conclusion of the present project. It was replaced by a much more efficient instrument, the SOPHIE spectrograph (Bouchy et al. 2006), which can attain a precision of 3-4 m s−1 in the radial- velocity measurements. A French-Swiss consortium (Loeillet et al. 2006), led by Fran¸cois

34 3.3. Calibrations using the CCF profile

Figure 3.2: Dependence of the CORALIE CCF profile on the projected rotation velocity of the star (left), and on its metallicity and effective temperature (right). Figure from Santos et al. (2002).

Bouchy, was created aiming to explore the high capabilities of the new instrument in a very large scientific programme, composed of five subprogrammes. With ELODIE, we observed at least once 854 objects from our sample of 1061 stars, and the new instrument continue observing 1) the stars with no observations done with ELODIE, 2) the stars with only 1 or 2 measurements and having [Fe/H] 0.1, and 3) the stars with significant ≥− radial-velocity variation but for which no orbital solution could be determined with the ELODIE observations. More than about 400 stars fulfil these criteria.

3.3 Calibrations using the CCF profile

The CCF represents a averaged profile of the individual spectral lines. This attribute of the CCF technique has important and very useful consequences. Since the spectral lines reflect what is going on in the stellar atmosphere, some astrophysical properties of the observed star can be derived from one single spectrum just by computing the CCF, and no high signal-to-noise spectrum is required to obtain good results. For solar-type stars, it is possible to have a good estimate of the projected rotation velocity of the star by measuring the excess in width of CCF profile in comparison to the width measured for non-rotating stars. Likewise, the equivalent width of the profile gives an estimate of the stellar metallicity with a quite good accuracy since, for a given effective temperature, the surface of the spectral lines strongly depends on the stellar metal content. Figure 3.2 shows the dependence of the CORALIE CCF profile on the projected rotation velocity, metallicity, and effective temperature.

35 Chapter 3. Sample selection and observations

Let A be the amplitude and σ the width of the CCF profile. Its equivalent width W , frequently called surface, is defined as:

W = √2πσA (km s−1) (3.1)

The surface of the CCF was proposed by Mayor (1980) to be used as a metal abundance estimator, given its correlation with the metallicity of the star. Benz & Mayor (1981) developed a method in which the width of the CCF profile gives information about the projected rotation velocity of the star (the v sin i parameter). Santos et al. (2002) recalled the main idea of these methods to construct a calibration using the CCF’s surface and width of the CORALIE spectrograph. In this work, similar calibrations were used, but using the width and surface of the ELODIE CCF.

3.3.1 Calibration of the projected rotation velocity

For stars having small rotation rate, the profile of the CCF calculated from their spectra can be approximated by a Gaussian function. In such a condition, the observed width of the CCF can be written in terms of its intrinsic width (the one expected for a non-rotating star) added quadratically to the broadening caused by the stellar rotation, which is assessed by the projected rotation velocity v sin i. The relation between v sin i and the observed width σ of the CCF can thus be written as:

v sin i =K σ2 σ2 (3.2) − 0 * where K is a constant, and σ0 is the mean intrinsic width of the CCF, which depends on the instrument profile, and on other broadening sources affecting the spectral line profile. I notice that this equation is valid for small values of v sin i (< 20 m s−1), and represents the excess in width of the line profile assuming that, for solar-type stars, the broadening is mostly caused by the stellar rotation velocity. Queloz et al. (1998) determined the constant K (1.9 0.1) using a sample of non-rotating dwarf stars observed with ELODIE. They ± constructed the following calibration of σ as a function of the colour index (B V ): 0 − σ =0.27(B V )2 +4.51 (km s−1) (3.3) 0 − where (B V ) represents the dependence on temperature. We used this calibration in − this work to estimate the v sin i of the sample stars. It was constructed for (B V ) in the − interval 0.6 (B V ) 1.4, but small extrapolations are allowed. ≤ − ≤

3.3.2 Calibration of the metal abundance

On his PhD thesis, Dominique Naef established a relation between the width of the ELODIE CCF and that of the CORALIE CCF, making possible to determine the metallic- ity of stars observed with ELODIE but using the calibration done by Santos et al. (2002)

36 3.3. Calibrations using the CCF profile

Figure 3.3: Left: Normalised surface of the CCF as a function of the maxcpp parameter for ELODIE stars having at least 4 spectra. The systematic effect for low quality spectra can be observed. W represents the mean value of the surface of the CCF for each ( ) star. Right: Surface of the ELODIE CCF as a function of (B V ) for the 43 stars with −− precise spectroscopic metallicity (those with maxcpp > 2500 e ). The calibration given by Equation 3.4 for [Fe/H] = 0 (dashed line), +0.25 and 0.25 dex (dotted lines) are also − plotted. Symbols represent different ranges of metallicity: red triangles: [Fe/H] 0.15; ≤− blue circles: 0.15 < [Fe/H] 0.15; magenta squares: [Fe/H] > 0.15. − ≤ for stars observed with CORALIE. After that, when precise spectroscopic metallicities were available for a large number of stars measured with ELODIE, he was able to con- struct a direct calibration of metallicity as a function of (B V ) and the surface of the − ELODIE CCF. Firstly, only stars having at least four measurements were selected (more than 1500 spectra of 323 stars meet this criterion). The minimum S/N required was de- termined taken into account a systematic effect observed in the surface of the ELODIE CCF: the smaller the spectral quality, the smaller the surface of the CCF. This effect, due to contamination by the thorium lamp, can be observed in Figure 3.3 (left panel), in which the normalised surface (W/ W ) is plotted against the spectral quality. Instead of ( ) using the ratio S/N, the spectral quality was checked by means of the maxcpp parameter, which represents the flux (in electrons) in the continuum of the CCF before normalisation. The plot shows that a minimum maxcpp of 2500 e− should be considered, which roughly corresponds to spectra with minimum S/N = 70 at λ5500 A.˚

A total of 43 stars having a precise spectroscopic determination of the metallicity were then selected among the high quality spectra (those having maxcpp > 2500 e−). The spectroscopic metallicities were determined by Nuno Santos also using high quality spectra. The calibration of metallicity as a function of the surface of the CCF and (B V ) −

37 Chapter 3. Sample selection and observations

Figure 3.4: Metallicity dependence on S/N for one of the programme stars (HD 185269). Comparing the three panels we observe that the systematic effect for low S/N comes essentially from the amplitude of the CCF. using these 43 ELODIE stars is represented in the right panel of Figure 3.3 and given by the following relation:

W B V B V 2 [Fe/H] = 4.055 log +0.648 0.768 − +0.126 − (3.4) 1.744 − 0.645 0.645 + ! " ! " ! " , which is valid in the intervals 0.38 [Fe/H] 0.33 and 0.488 (B V ) 1.024. A ≤ ≤ ≤ − ≤ superior limit of 0.9 in (B V ) seems to be better since we have only two isolated points − with (B V ) greater than such a value. In any case, most of the target stars have (B V ) − − between 0.5 and 0.9. The rms around the calibration is 0.064 dex.

3.3.3 The dependence on S/N

The systematic effect just described has a direct impact on the metallicity determined for stars with low quality spectra. Left panel of Figure 3.4 shows the dependence of the metallicity on the S/N ratio for one of the programme stars. Since most of the target stars have spectra with S/N between 20 and 70 (see Figure 3.7), the computed metallicities for such stars would be in fact underestimated if we used the proposed calibration. A correc- tion had thus to be applied taking the systematic effect into account. To determine this correction, we selected in the ELODIE database all stars having at least 4 measurements and, for each star, we calculated the normalised value for the surface of the CCF. We then fitted a curve to the data points in the W/ W maxcpp diagram (see Figure 3.5). This ( )× curve represents the correction applied to all surface values of the observed stars, given their maxcpp value, and is set by the following relation:

1 20.29 W = +0.0042 ( 0.007) (3.5) corr 1 + 0.0195 exp( maxcpp/349) − maxcpp ± − 38 3.3. Calibrations using the CCF profile

Figure 3.5: Left: Curve fitted to the normalized surface of the CCF as a function of the maxcpp parameter (given by Equation 3.5), representing the correction applied to the surface values of the observed stars. Right: Surface of the CCF after the correction.

Comparing the three panels of Figure 3.4, and remembering that the metallicity is correlated with the surface, which in turn is a function of the CCF’s width and amplitude (Equation 3.1), we can notice that the effect essentially comes from the amplitude of the CCF. Nevertheless, we applied the corrections directly to the surface of the CCF. There- fore, to estimate the metallicity of our target stars that is unaffected by the systematic effect occurring when the spectral quality is not high enough, the surface W of the CCF of each star was replaced in Equation 3.4 by the cleaned value Wclean, corrected using Equation 3.5 as follows: W −1 Wclean = (km s ) (3.6) Wcorr

39 Chapter 3. Sample selection and observations

Figure 3.6: Left: Distribution of the number of measurements for each observed star having [Fe/H] 0.1 and v sin i 6 km s−1. A total of 1359 spectra were gathered using ELODIE ≥ ≤ for 221 of such stars in the present programme. The last bin includes the stars with more than 20 measurements. Right: Distribution of the radial-velocity dispersion for 209 stars having at least 3 measurements made with ELODIE. The last bin represents the stars with RV dispersion larger than 100 m s−1.

3.4 Global view of the observations

Here I present the distribution of some parameters related to the radial-velocity mea- surements made for the programme stars using the ELODIE spectrograph: the number of measurements per star, the radial-velocity dispersion, the S/N ratio of the collected spectra, the estimated photon-noise error of each measurement, the projected rotation velocity, and the metallicity.

The left panel of Figure 3.6 presents the distribution of the number of radial-velocity measurements made for the metal-rich ([Fe/H] 0.1) and low rotation (v sin i 6 km s−1) ≥ ≤ stars observed with ELODIE in the context of this programme. We collected a total of 1359 spectra of 221 of such stars. I did not include in this distribution neither the stars poorer than [Fe/H] = 0.1 nor those with v sin i larger than 6 km s−1 since, in general, they were not monitored according to the adopted observational strategy (see Section 3.2). Notice in this figure that most stars have at least 3 measurements. Those having only 1 or 2 spectra need more measurements and should be observed with the SOPHIE spectrograph before being classified as radial-velocity constants, or possible candidates to harbour a planetary companion (in Chapter 6, I present our criteria used to select a subsample of stars to be monitored with SOPHIE).

40 3.4. Global view of the observations

Figure 3.7: Left: Distribution of S/N ratios at the echelle order 47 (λλ5520 5580) for − 2341 spectra of the 854 observed stars. Right: Photon-noise errors as a function of S/N at the order 47 for the 2341 measurements. Circles indicate points having widths larger than 8 km s−1, which is the case of fast-rotating stars.

On the right panel of Figure 3.6 I show the distribution of the radial-velocity dispersion for the observed stars for which we made at least 3 measurements. A total of 209 stars are represented in the histogram.

The left panel of Figure 3.7 shows the distribution of the S/N ratios measured at the echelle order 47 of the spectra at the moment of the automated reduction. This order contains the wavelength range between 5520 and 5580 A.˚ A total of 2341 measurements are represented in the figure, and correspond to all spectra collected in the course of the programme. The right panel of the same figure presents the estimated photon-noise error as a function of the S/N ratio for each one of the 2341 measurements. They were derived using Equation 3.9 (see below). The width σ of the CCF have an important impact on the estimate of the photon-noise errors: when the CCF is large (normally due to the stellar rotation), its profile deviates from a Gaussian and the fitted function is less precise, inducing larger errors. Some measurements for which the CCF have a very large profile (σ > 8 km s−1) are indicated in the figure. Spectra having S/N smaller than 20 were normally not taken into account given their large photon-noise errors.

Figure 3.8 shows the distribution of both the projected rotation velocity v sin i and metallicity [Fe/H] determined for all the 854 measured using the ELODIE spectrograph in the course of our observations.

41 Chapter 3. Sample selection and observations

Figure 3.8: Distributions of projected rotation velocity v sin i (left) and metallicity [Fe/H] (right) for the 854 stars observed with ELODIE in the context of our programme. The last bin in the v sin i distribution includes all stars having v sin i>10 km s−1.

3.4.1 Photon-noise errors

Baranne et al. (1996) estimated the uncertainties due to photon noise for the ELODIE spectrograph by doing Monte Carlo simulations, and found the following relation:

C(T )(1 + 0.2w) ε = eff (km/s) (3.7) ph 3A S/N ( ) where C(Teff ) is a constant that depends on the spectral type of the star and on the numerical mask used, w and A are respectively the width (FWHM) and the relative depth or amplitude of the CCF, and S/N represents the global signal-to-noise ratio of the ( ) spectrum. In the case of ELODIE, C(Teff ) = 0.042 for the masks cleaned from telluric lines (Naef 2004). The global S/N can be estimated using the relation: maxcpp S/N = (3.8) 2 ( ) (maxcpp + NpixL ) where maxcpp (measured in number- of electrons) is the parameter defined in Section 3.3, Npix (measured in pixels) is the height of the echelle orders in the CCD, and L is the − readout noise of the CCD. For ELODIE, Npix = 3 pixels and L = 8 e . Assuming that √ FWMH = 2 2 ln 2 σ ∼= 2.355 σ , we can rewrite Equation 3.9 as:

2 0.042 (1 + 0.2 2.355 σ) (maxcpp + NpixL ) εph = · (km/s) (3.9) 3 A maxcpp-

42 3.4. Global view of the observations

where σ is the standard deviation of the Gaussian function fitted to the CCF. This is the relation we used to determine the photon-noise errors of our radial-velocity measurements made using the ELODIE spectrograph.

Before discussing in the next chapter some particular cases concerning our results, I present in Table 3.2 a global view of the stars observed in the context of our programme. Column 2 lists the number of ELODIE measurements (N) for stars observed at least 3 times, except some particular cases for which some observations were also made using the SOPHIE or HARPS spectrographs. Columns 3 and 4 list, respectively, the stellar metallicity and projected rotation velocity calculated from the calibrations of Section 3.3. Columns 5, 6 and 7 show, respectively, a mean value for the radial velocities RV , their ( ) amplitude ∆ RV, and dispersion σ , which were calculated only when N 2. Column 8 RV ≥ lists the ratio between external and internal dispersions E/I (see definition in Section 4.5), whereas column 9 shows the range ∆ JD between first and last measurements. The last column presents some comments in order to better explain some particular cases.

43 Chapter 3. Sample selection and observations

Table 3.2: Global results for stars observed with ELODIE in the context of our main programme. I list the number of ELODIE measurements N (only those having N ≥ 3, except some particular cases), metallicity and projected rotation velocity from the calibrations of Section 3.3, a mean value for the radial velocities RV , their amplitude ( ) of variation ∆ RV and dispersion σRV , the external-to-internal dispersion ratio E/I, and the range of the measurements ∆ JD. The last column presents some comments in order to better explain some particular cases (see legend at the end of this table).

v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − −1 E/I Remarks [km s 1] [km s−1] [km s−1] [m s ] [days] BD +394559 4 0.24 2.9 −64.226 0.006 2.9 0.3 243 BD +481523 1 0.07 0.1 −66.497 - - - - RV var, Bin BD +600583 4 0.31 2.5 −1.330 0.042 20.7 1.8 6 BD +630438 3 0.11 9.1 18.283 0.025 14.2 0.8 57 HD 10790 0 ------RV var, Bin HD 100796 28 0.12 4.2 42.645 0.178 40.9 1.8 507 RV var, Pec HD 102195 0 ------Pl, H HD 104289 7 0.07 4.9 −20.447 2.257 1078.7 70.2 421 RV var, Bin HD 105279 3 0.27 3.7 −13.192 0.012 6.2 0.3 91 HD 105585 3 0.29 3.7 0.743 0.035 20.2 1.2 91 HD 105844 5 0.20 0.0 0.492 0.022 9.0 1.1 451 HD 106888 2 0.14 5.8 −3.745 0.010 7.1 0.3 0 RV var, Bin HD 108076 3 −1.06 0.2 −2.441 0.057 31.8 1.3 571 HD 108575 24 0.16 5.6 −1.643 0.172 40.8 3.3 713 RV var, Act HD 108942 4 0.20 2.1 −10.963 0.026 12.2 1.4 96 HD 109552 3 0.19 3.7 24.676 0.034 17.2 1.7 31 HD 110869 3 0.13 1.4 −9.144 0.023 11.7 1.3 31 HD 110950 3 0.15 3.6 6.276 0.041 21.0 0.9 381 HD 111395 5 0.15 2.6 −8.998 0.040 14.8 3.9 85 HD 11168 3 0.08 3.0 13.422 0.030 17.0 0.7 380 HD 112001 3 0.19 1.8 −12.203 0.037 20.6 2.0 361 HD 11271 3 0.21 3.3 9.313 0.039 20.6 2.0 348 HD 113998 9 0.16 1.9 14.629 0.533 160.2 4.6 438 HD 114036 3 0.27 1.6 0.282 0.032 17.0 2.4 274 HD 114285 3 0.07 4.2 −29.825 0.035 17.7 1.1 2 HD 115954 4 0.29 4.5 −14.799 0.059 25.8 1.8 278 HD 116091 4 0.30 3.3 7.902 0.055 26.1 2.1 314 HD 11638 14 0.12 2.5 11.436 0.253 61.9 4.5 501 HD 117655 3 0.37 1.8 −23.018 0.029 15.5 2.2 409 HD 118203 56 0.15 4.7 −29.310 0.544 159.9 12.4 750 Pl, E HD 121763 3 0.26 1.5 −39.188 0.025 14.2 1.3 2 HD 121979 25 0.24 3.5 −15.355 0.096 24.0 1.7 449 RV var, Act HD 122727 19 0.26 3.9 8.722 0.075 18.7 1.4 341 HD 123300 3 0.10 3.2 −32.451 0.006 3.2 0.3 658 HD 124330 3 0.28 1.7 −30.747 0.020 11.3 0.7 385 HD 125193 17 0.28 3.4 −31.697 1.489 491.8 58.5 742 RV var, Bin

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v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − − −1 E/I Remarks [km s 1] [km s−1] [km s 1] [m s ] [days] HD 126945 6 0.34 2.7 −12.033 0.063 22.0 1.9 422 HD 129171 7 0.18 2.6 −11.873 0.034 10.9 1.2 484 HD 130215 3 0.07 1.6 −20.519 0.010 5.0 0.9 23 HD 131526 6 0.23 3.4 −45.699 0.038 16.6 2.6 58 HD 132406 17 0.18 1.7 −37.797 0.201 57.5 4.7 747 Pl, E+S HD 132505 3 0.26 3.7 −15.498 0.007 3.8 0.5 60 HD 134353 4 0.03 1.4 −24.058 0.088 40.1 2.8 7 HD 135633A 3 0.30 4.6 −30.144 0.035 19.7 1.2 108 HD 136064 3 0.06 5.0 −48.101 0.014 8.1 1.2 402 HD 136902 5 0.18 1.5 −41.910 0.033 15.1 1.0 9 HD 13836 3 0.16 1.9 1.211 0.014 7.0 0.9 54 HD 13997 5 0.23 0.5 −20.891 0.014 6.1 0.6 329 HD 142925 3 0.23 3.1 −13.443 0.012 6.0 0.4 69 HD 144302 3 0.23 4.4 −39.103 0.009 4.6 0.4 359 HD 14651 1 0.04 0.0 53.640 - - - - RV var, Bin HD 147187 22 0.18 1.4 −22.260 0.058 16.1 1.9 746 RV var, Pec HD 147887 3 0.09 2.8 6.786 0.024 12.1 0.5 258 HD 148164 4 0.13 5.4 −64.623 0.121 59.1 2.2 5 HD 149143 14 0.20 3.8 12.147 0.352 113.3 8.4 351 Pl, E+O HD 149222 5 0.18 2.3 −33.923 0.095 36.2 2.4 388 HD 150633 3 0.12 1.8 −76.174 0.035 18.3 1.7 16 HD 151501 26 0.29 4.5 16.403 0.103 29.4 2.0 423 HD 15210 13 0.06 4.8 8.532 0.159 43.9 2.5 98 HD 15292 12 −0.00 1.2 −37.887 2.171 716.8 67.7 150 RV var, Bin HD 15397 8 0.11 4.9 −0.103 0.041 16.0 1.2 61 HD 154160 4 0.30 1.3 −52.870 0.008 3.3 0.7 421 HD 155358 0 ------Pl, O HD 155456 8 0.14 0.4 −59.677 0.067 25.1 2.2 350 HD 156279 3 0.08 0.9 −20.523 0.721 409.8 36.5 447 HD 157102 4 0.32 0.7 −42.692 0.035 14.9 2.6 327 HD 158332 14 0.20 0.2 −24.216 0.040 11.3 1.4 804 HD 15851 5 0.13 2.2 10.041 0.050 18.5 1.6 94 HD 15942 3 0.40 2.3 34.190 0.019 9.5 1.8 347 HD 160013 4 0.19 0.8 1.382 0.040 18.5 1.2 312 HD 161479 4 0.19 3.9 −33.508 10.75 5864.2 612.0 336 RV var, Bin HD 16175 3 0.29 4.4 21.779 0.005 2.9 0.5 52 HD 163589 18 0.19 1.3 −35.978 0.073 18.2 2.2 365 RV var, Pec HD 168960 4 0.11 3.5 −10.107 0.026 11.5 1.3 15 HD 169925 3 0.25 4.7 −21.835 0.014 7.4 1.3 336 HD 17156 2 0.04 4.5 −3.127 0.040 28.3 1.1 12 Pl, S+O HD 172085 3 0.25 4.0 −34.340 0.024 12.7 1.1 72 HD 172669 3 0.11 2.5 −6.708 0.006 3.2 0.4 331 HD 173700 7 0.25 6.0 −17.811 0.039 13.7 1.3 268 HD 175441 3 0.20 6.1 −21.585 0.008 4.0 0.4 409 HD 180161 6 0.15 2.3 −27.144 0.048 17.8 3.3 379

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45 Chapter 3. Sample selection and observations

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v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − − −1 E/I Remarks [km s 1] [km s−1] [km s 1] [m s ] [days] HD 180263 6 0.12 1.8 −58.023 0.064 22.6 2.2 35 HD 180502 7 0.15 3.6 −6.907 0.053 21.1 1.6 745 HD 183263 7 0.17 2.8 −50.408 0.158 57.7 4.0 744 HD 185239 6 −0.14 5.7 11.252 0.067 23.1 1.5 40 HD 185269 44 0.15 5.5 0.615 0.214 60.0 6.6 408 Pl, E+O HD 186843 6 0.06 3.0 −30.399 0.187 68.2 2.9 372 HD 18702 3 0.13 0.2 71.249 0.004 2.0 0.3 481 HD 187882 3 −0.51 2.3 −55.390 0.067 33.7 1.9 323 HD 189733 26 −0.07 3.5 −2.356 0.427 102.1 14.4 349 Pl, E HD 190470 3 0.07 1.5 −7.260 0.001 0.6 0.1 22 HD 191806 6 0.31 3.2 −15.481 0.215 102.6 9.5 685 HD 193215 8 0.10 2.7 −1.085 0.080 23.6 2.1 655 HD 197207 6 0.12 2.0 −56.621 0.033 11.7 1.0 54 HD 197488 3 0.14 4.0 9.163 0.024 13.3 1.2 339 HD 19902 3 0.18 2.0 27.308 0.010 5.8 0.9 55 HD 199100 5 0.34 0.9 −20.973 2.841 1486.6 325.8 54 RV var, Bin HD 200078 3 0.24 1.5 −60.250 0.036 18.1 1.7 24 HD 200254 3 0.27 2.2 −18.870 0.024 12.1 1.5 21 HD 200560 3 0.06 3.6 −14.263 0.014 7.2 1.5 350 HD 201219 3 0.19 2.8 4.834 0.034 17.8 1.1 11 HD 20278 3 0.14 4.6 46.287 0.024 12.5 1.0 61 HD 203698 4 0.13 5.1 −16.482 0.026 11.9 1.4 719 HD 204906 3 0.05 7.5 13.044 0.043 24.8 1.0 100 HD 205351 4 0.05 3.1 −18.042 0.050 22.6 1.3 681 HD 205702 7 0.12 4.5 −13.491 0.047 15.9 1.1 57 HD 206772 6 0.11 2.3 −10.268 0.059 22.5 2.4 659 HD 20678 3 0.11 3.4 35.224 0.034 18.4 1.9 441 HD 207839 20 0.20 1.9 −29.908 0.061 17.7 2.1 743 HD 208552A 9 0.07 4.2 −19.243 0.100 31.7 1.2 10 HD 208863 3 0.18 3.5 11.746 0.011 6.1 0.5 38 HD 210144 4 0.11 0.2 −33.370 0.008 3.4 0.6 720 HD 211275 5 0.20 3.6 −45.833 0.025 9.9 0.9 367 HD 211403 10 0.27 17.6 −9.269 0.444 148.8 2.5 532 HD 211681 12 0.26 2.4 −40.124 0.266 81.8 6.6 745 HD 212585 14 0.09 3.5 −15.820 1.395 407.1 40.0 721 RV var, Bin HD 212733 8 0.06 0.1 10.510 7.216 3248.0 380.7 603 RV var, Bin HD 214823 5 −0.03 5.4 −44.738 0.355 146.8 7.3 374 HD 216191 3 −0.02 0.0 18.953 0.042 24.0 1.1 307 HD 216320 8 −0.13 2.8 −12.582 13.842 4243.1 293.0 535 RV var, Bin HD 21742 4 0.14 0.8 −36.173 0.010 4.1 0.7 88 HD 21774 3 0.30 1.5 −3.150 0.016 8.1 1.1 59 HD 219428 3 0.10 4.8 −6.448 0.036 20.0 1.3 661 HD 219828 3 0.18 2.6 −24.126 0.049 24.6 2.8 52 Pl, H HD 220221 5 −0.09 2.2 −13.584 0.022 8.0 0.9 11 HD 221585 10 0.20 0.8 6.347 0.092 31.5 5.1 655

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v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − − −1 E/I Remarks [km s 1] [km s−1] [km s 1] [m s ] [days] HD 221627 16 0.21 2.8 −5.728 0.092 21.4 2.9 113 HD 224508 19 0.09 2.7 28.817 0.138 33.5 2.1 372 HD 2330 4 0.25 3.9 −8.677 0.003 1.7 0.2 57 HD 24053 12 0.13 2.4 4.385 0.060 18.9 1.7 493 RV var, Act HD 26756 3 0.19 4.9 38.293 0.048 24.0 1.7 123 HD 27282 3 0.21 5.0 37.987 0.018 9.6 1.0 56 HD 27495 3 0.19 2.7 −14.273 0.022 11.1 1.8 121 HD 276618 4 0.22 0.8 18.753 0.002 1.2 0.1 270 HD 27969 3 0.17 3.2 47.164 0.012 6.7 0.6 66 HD 28099 34 0.22 3.4 38.580 0.127 28.2 2.0 503 RV var, Act HD 28635 3 0.21 3.9 40.616 0.038 19.0 0.9 6 HD 29037 1 0.02 7.5 23.530 - - - - RV var, Bin HD 29836 21 0.28 2.1 13.411 0.253 58.5 6.7 421 RV var, Pec HD 29862 3 0.22 1.7 20.138 0.025 12.5 1.9 3 HD 30246 11 0.11 4.2 41.850 0.130 41.2 2.1 56 HD 30376 3 0.13 0.5 −45.554 0.008 4.2 0.4 35 HD 30572 11 0.33 6.2 32.445 0.112 29.2 1.7 62 RV var, Act HD 30589 9 0.19 5.2 41.517 0.116 36.0 1.9 23 RV var, Pec HD 30925 3 0.17 3.6 −13.577 0.035 19.4 1.6 53 HD 31000 25 0.17 4.7 −4.765 0.132 30.3 2.7 476 RV var, Act HD 31338 4 0.25 0.0 30.322 0.067 31.0 2.4 327 HD 3141 10 0.06 0.1 0.593 0.043 14.5 1.9 125 RV var, Pec HD 31452 3 0.14 0.2 15.051 0.040 21.4 1.7 31 HD 31781 3 0.13 2.9 16.707 0.020 10.8 0.8 116 HD 3251 3 0.21 2.7 −15.447 0.039 19.5 2.0 290 HD 34031 4 0.13 2.4 23.474 9.840 4204.7 463.2 543 RV var, Bin HD 34887 4 0.28 0.7 −25.586 0.020 8.7 1.2 22 HD 35768 3 0.08 5.3 42.301 0.010 5.0 0.2 325 HD 36130 3 0.13 1.8 −62.473 0.016 8.5 0.7 7 HD 36248 5 0.14 3.8 2.903 0.051 24.7 1.7 49 HD 3758 3 0.14 2.6 6.497 0.035 17.5 1.4 34 HD 40330 6 0.07 2.7 56.706 0.076 28.1 1.3 271 HD 4075 3 0.29 1.6 −9.364 0.024 13.9 2.4 42 HD 43383 7 0.06 3.0 12.072 2.357 883.1 47.6 687 HD 43691 24 0.22 4.7 −28.968 0.283 84.6 5.8 538 Pl, E+S HD 45652 19 0.15 0.2 −5.113 0.169 50.5 4.5 151 Pl, E+C HD 4635 3 0.07 1.5 −31.563 0.013 6.7 1.0 318 HD 4649 19 0.03 2.9 −26.412 0.201 56.6 3.2 182 HD 4868 3 0.28 3.1 −5.872 0.006 3.1 0.4 30 HD 49178 1 0.06 1.1 48.870 - - - - RV var, Bin HD 49385 3 0.18 3.8 33.184 0.026 13.1 1.0 87 HD 5470 14 0.21 2.1 −3.796 0.468 176.1 13.4 162 HD 54807 4 0.09 2.6 −1.989 0.041 17.4 0.7 2 HD 58900 3 0.18 7.6 49.304 0.058 30.4 1.1 32 HD 59062 3 0.23 1.1 45.876 0.042 21.9 1.6 30

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47 Chapter 3. Sample selection and observations

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v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − − −1 E/I Remarks [km s 1] [km s−1] [km s 1] [m s ] [days] HD 60521 3 0.19 3.1 29.427 0.010 5.0 0.4 266 HD 62923 15 0.20 1.9 6.130 4.107 1796.5 185.7 57 RV var, Bin HD 6512 3 0.14 2.1 10.310 0.036 18.0 2.1 116 HD 6660 3 −0.11 2.3 6.556 0.010 5.5 0.7 116 HD 6697 3 0.24 0.8 −24.714 0.017 9.1 1.3 117 HD 67346 3 0.34 4.0 26.770 0.011 5.9 0.6 88 HD 68255A 26 0.22 4.5 −5.633 0.162 47.9 7.7 421 HD 69056 3 0.11 0.0 20.520 0.038 19.1 2.2 334 HD 691 16 0.28 5.2 −2.867 0.108 24.6 2.1 541 RV var, Act HD 72616 7 0.36 2.0 23.871 0.064 24.2 1.4 330 HD 7352 27 0.15 3.9 −17.273 0.110 32.7 1.9 501 RV var, Pec HD 7355 10 0.39 4.2 −6.949 0.382 106.9 11.0 62 HD 75528 31 0.09 3.1 45.651 0.144 29.7 3.7 802 RV var, Pec HD 75576 3 0.10 4.3 −12.354 0.069 35.4 1.9 5 HD 75898 21 0.21 4.0 21.792 0.141 39.5 2.4 477 Pl, E+O HD 76025 4 0.19 3.4 29.050 0.038 16.1 0.9 363 HD 76539 3 0.19 2.0 15.274 0.029 15.7 1.7 3 HD 76974 3 0.19 3.6 −40.346 0.010 5.3 0.6 4 HD 78277 6 0.35 3.2 2.675 0.061 22.1 1.6 57 HD 78536 1 0.07 4.5 3.417 - - - - RV var, Bin HD 80869 23 0.13 2.3 −18.212 0.188 50.5 2.8 479 RV var, Pec HD 81505 7 0.30 1.7 17.104 0.100 41.4 3.5 115 HD 82443 16 0.11 5.8 8.391 0.240 62.4 6.2 559 RV var, Act HD 82881 3 0.10 2.4 24.035 0.019 9.7 0.7 691 HD 82885 3 0.35 2.4 14.433 0.023 13.0 3.3 23 HD 82939 3 0.10 1.7 −0.051 0.019 10.2 0.9 362 HD 84869 3 0.14 5.0 −24.899 0.045 22.8 0.9 3 HD 85362 3 0.14 1.6 1.196 0.036 18.3 1.7 240 HD 86460 22 −0.21 1.2 3.127 0.701 137.0 7.7 449 RV var, Pec HD 8745 4 0.12 3.8 −8.335 0.053 25.1 2.3 52 HD 87680 3 −0.12 1.7 −26.926 0.042 21.0 1.2 70 HD 87912 3 0.24 2.1 24.780 0.032 17.2 2.0 82 HD 88232 4 0.30 1.2 30.435 0.046 22.9 1.1 418 HD 88402 3 0.13 3.1 34.473 0.029 14.5 1.5 270 HD 89010 9 0.11 2.6 −35.463 2.149 664.9 100.9 384 RV var, Bin HD 90681 6 0.25 3.9 4.151 0.060 20.5 1.9 294 HD 9070 3 0.35 1.5 11.677 0.006 3.0 0.5 292 HD 9081 3 0.26 2.4 28.750 0.025 12.7 1.3 58 HD 91163 10 0.06 4.6 −21.367 0.073 22.7 1.4 410 HD 91332 4 0.32 3.9 −43.599 0.058 26.9 1.1 324 HD 91988 7 0.13 2.7 43.950 0.086 33.6 2.2 294 HD 94667 3 0.26 3.1 −89.344 0.034 17.8 1.2 140 HD 94835 14 0.12 2.5 8.208 0.383 96.2 6.1 404 HD 94861 3 0.18 6.0 −5.551 0.021 10.5 0.4 2 HD 96853 5 0.16 0.6 0.744 0.047 19.4 1.9 660

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48 3.4. Global view of the observations

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v sin i " RV # ∆ RV σRV ∆ JD Star ID N [Fe/H] − − −1 E/I Remarks [km s 1] [km s−1] [km s 1] [m s ] [days] HD 96937 3 0.12 0.6 10.087 0.006 3.1 0.3 115 HD 97140 8 0.25 4.0 −34.370 0.023 7.0 0.7 447 HD 97657 7 0.12 5.1 −7.567 0.124 43.5 1.8 144

Pl: stars with detected planets. The instruments used are indicated: E: ELODIE, C: CORALIE, H: HARPS, S: SOPHIE, O: Others. See Table 4.2 for a complete list of such host stars. RV var: stars with significant radial velocity variation but with no planet detected. They are presented in Section 4.5. Act: chromospheric active and/or photometric variable stars. They are discussed in Subsection 4.5.1. Pec: Some peculiar cases, for which no information is available in the literature to explain the observed radial-velocity variation. They are presented in Subsection 4.5.2. Bin: stars revealed to be a member of a binary system, having either a faint low mass star or a brown dwarf as a companion. They are discussed in Subsection 4.5.3.

49 Chapter 3. Sample selection and observations

50 Chapter 4

Main results

Contents

4.1 The discovery of a hot Jupiter transiting HD 189733 ...... 51 4.2 Other publications concerning the HD 189733 system . . . . . 59 4.3 Other hot Jupiters discovered with our programme ...... 62 4.4 Additional discoveries ...... 77 4.5 Radial-velocity variable stars with no detected planet . . . . . 101

4.1 The discovery of a hot Jupiter transiting HD 189733

HD 189733, one of the selected target stars from our programme, is a K dwarf star of visual magnitude V =7.67, situated in the northern hemisphere in a distance of 19.3 pc from the Sun. The low metallicity of this star ([Fe/H] = 0.03) does not fulfil the initial − observing strategy. However, since the first observation did not have the spectral quality needed for an accurate metallicity estimate, we gathered a second spectra of HD 189733. This revealed a significant radial-velocity variation, which led us to follow this object with more observations. After the first 8 radial-velocity measurements, we were able to determine an orbital solution for the system: a 1.3 Jupiter-mass planet orbiting the main star with a very short period of 2.2 days. A very hot Jupiter had been discovered. ∼ Given the high probability of a transit event (11.3%), we decided to conduct spectro- scopic and photometric observations of HD 189733 during the estimated transit time. A dimming in the photometric light curve and a Rossiter-McLaughlin effect manifested in the radial velocities both revealed that HD 189733 b indeed transit the disk of its parent star as viewed from Earth. This provided the determination of some physical parameters

51 Chapter 4. Main results

of the planetary companion: a mass of 1.15 0.04 MJup, a radius of 1.26 0.03 RJup and − ± ± a mean density of 750 80 kg m 3. ± The planet transiting HD 189733 was firstly published by Bouchy et al. (2005), and the paper of announcement is presented in the following subsection. It is the ninth tran- siting extrasolar planet found, and the third one transiting a bright star. A relative large planet transiting a relative small star (0.76 R!), the proximity of the system, and the large apparent magnitude of the star, all together show the very relevant nature of this system. For these reasons, many other works have been published since then, discussing and improving its main physical properties. In next section, I present some of these works and their main results.

4.1.1 Paper: A very hot Jupiter transiting the bright star HD189733

F. Bouchy, S. Udry, M. Mayor, C. Moutou, F. Pont, N. Iribarne, R. Da Silva, S. Ilovaisky, D. Queloz, N. C. Santos, D. S´egransan, and S. Zucker

Bouchy et al. 2005, Astronomy & Astrophysics 444, L15-L19

1. Introduction

2. Observations

3. Stellar characteristics of HD 189733

4. Keplerian solution and spectroscopic transit

5. Photometric transit and characterisation of HD 189733 b

6. Summary and concluding remarks

52 4.1. The discovery of a hot Jupiter transiting HD 189733

A&A 444, L15–L19 (2005) Astronomy DOI: 10.1051/0004-6361:200500201 & !c ESO 2005 Astrophysics

ELODIE metallicity-biased search for transiting Hot Jupiters!

II. A very hot Jupiter transiting the bright K star HD 189733

F. Bouchy1,2, S. Udry3, M. Mayor3, C. Moutou1, F. Pont3, N. Iribarne2, R. Da Silva3, S. Ilovaisky2, D. Queloz3, N. C. Santos3,4, D. Ségransan3, and S. Zucker3,5

1 Laboratoire d’Astrophysique de Marseille, Traverse du Siphon, 13013 Marseille, France e-mail: [email protected] 2 Observatoire de Haute-Provence, 04870 St Michel l’Observatoire, France 3 Observatoire de Genève, 51 Ch. des Maillettes, 1290 Sauverny, Switzerland 4 Lisbon Observatory, Tapada da Ajuda, 1349-018 Lisboa, Portugal 5 Weizmann Institute of Science, PO Box 26, Rehovot 76100, Israel

Received 29 September 2005 / Accepted 14 October 2005

ABSTRACT

Context. Among the 160 known exoplanets, mainly detected in large radial-velocity surveys, only 8 have a characterization of their actual mass and radius thanks to the two complementary methods of detection: radial velocities and photometric transit. Aims. We started in March 2004 an exoplanet-search programme biased toward high-metallicity stars which are more frequently host extra-solar planets. This survey aims to detect close-in giant planets, which are most likely to transit their host star. Methods. For this programme, high-precision radial velocities are measured with the ELODIE fiber-fed spectrograph on the 1.93-m telescope, and high-precision photometry is obtained with the CCD Camera on the 1.20-m telescope, both at the Haute-Provence Observatory. Results. We report here the discovery of a new transiting hot Jupiter orbiting the star HD 189733. The planetary nature of this object is confirmed by the observation of both the spectroscopic and photometric transits. The exoplanet HD 189733 b, with an orbital period of 2.219 days, has one of the shortest orbital periods detected by radial velocities, and presents the largest photometric depth in the light curve (∼3%) observed to date. We estimate for the planet a mass of 1.15 ± 0.04 MJ and a radius of 1.26 ± 0.03 RJ. Considering that HD 189733 has the same visual magnitude as the well known exoplanet host star HD 209458, further ground-based and space-based follow-up observations are very promising and will permit a characterization of the atmosphere and exosphere of this giant exoplanet.

Key words. stars: individual: HD 189733 – planetary systems – techniques: radial velocities – techniques: photometry

1. Introduction (Konacki et al. 2003), OGLE-TR-113b and 132b (Bouchy et al. 2004), TrES-1 (Alonso et al. 2004), OGLE-TR-111b Within the last ten years more than 150 planetary systems have (Pont et al. 2004) and OGLE-TR-10b (Bouchy et al. 2005; been discovered, mostly by radial-velocity surveys. Although Konacki et al. 2005). Only 2 transiting exoplanets have they provide a great deal of information on the system orbital been discovered first by radial velocities and then had their parameters and on the host star properties, such surveys neither photometric transit measured: HD 209458 (Mazeh et al. 2000; yield the acurate mass of the planet (only m sin i) nor give any Charbonneau et al. 2000; Henry et al. 2000) and HD149026 information about its size. The observation of planetary transits (Sato et al. 2005). The host stars of the 5 OGLE planets are together with radial-velocity measurements yield, on the other unfortunately faint and complementary follow-up observations hand, the actual mass and planetary radius (and thus mean den- are difficult and time-consuming. We present in this letter a sity), providing constraints for planet interior models. new transiting hot Jupiter orbiting the bright (V = 7.7) and Within the past two years, 6 exoplanets have been dis- close (d = 19 pc) star HD 189733. covered first by the photometric observation of a transit in front of the stellar disk (e.g., Udalski et al. 2002) and then confirmed by spectroscopic follow-up: OGLE-TR-56b 2. Observations

! Based on observations collected with the ELODIE spectrograph HD 189733 belongs to our “ELODIE metallicity-biased search on the 1.93-m telescope and with the 1.20-m telescope (Observatoire for transiting Hot Jupiters” survey (Da Silva et al. 2005). This de Haute-Provence, France). survey started in March 2004 with the ELODIE spectrograph

53 Chapter 4. Main results

L16 F. Bouchy et al.: The transiting hot Jupiter HD 189733 b. II.

(Baranne et al. 1996) on the 1.93-m telescope at the Haute- Table 1. Parameters for the star HD 189733, for the Keplerian solution Provence Observatory (France). The main idea of the pro- and inferred for the planetary companion. gramme is to bias the target sample for high-metallicity stars which are more likely to host planets (Gonzales 1998; Santos Period [days] 2.219±0.0005 et al. 2001, 2005; Fischer & Valenti 2005). It already allowed Orbital eccentricity 0 [fixed] −1 the discovery of 2 hot Jupiters orbiting the stars HD 118203 Radial velocity semi-amplitude [ms ] 205 ± 6 −1 − ± and HD 149143 (the latter recently announced by Fischer et al. Systemic velocity [km s ] 2.361 0.003 O–C residuals [ms−1] 15 in prep.). The observational strategy of the survey is designed to pri- Transit epoch [JD-2 453 000] 629.3890 ± 0.0004 marily target hot Jupiters which are ideal candidates for follow- Radius ratio 0.172 ± 0.003 up photometric-transit searches. In practice, the first spectro- Impact parameter 0.71 ± 0.02 scopic measurement is made to estimate the metallicity by Inclination angle [◦] 85.3 ± 0.1 measuring the surface of the cross-correlation function of the ELODIE spectrum (Santos et al. 2002). Then, the star is se- Temperature [K] 5050±50 ± lected for further observations if the derived metallicity [Fe/H] log g 4.53 0.14 [Fe/H] −0.03 ± 0.04 is greater than 0.1 dex. The metallicity estimate requires a spec- v sin i [km s−1] 3.5 ± 1.0 trum with S/N ≥ 40. For HD 189733, this S/N ratio was Star mass [M(] 0.82 ± 0.03 only reached after two independent exposures which, more- Star radius [R(] 0.76 ± 0.01 over, revealed a large velocity variation. The object was there- Prot [days] ∼11 fore followed despite the fact that the measured metallicity was not larger than 0.1. Typical exposure times were between Orbital semi-major axis [AU] 0.0313 ± 0.0004 15 and 25 min, corresponding to photon-noise uncertainties of Planet mass [MJ] 1.15 ± 0.04 −1 ∼5–7 ms . Planet radius [RJ] 1.26 ± 0.03 Planet density [g cm−3] 0.75 ± 0.08 A first set of 8 radial velocities of HD 189733, showing a large-amplitude variation, allowed us to easily constrain a cir- cular orbital solution with a very short period (2.2 days) for the companion. With such a short period, the probability that the obtained with the CORALIE spectrograph gives Teff = 5050 ± companion crosses the stellar disk is quite high (∼1/8). We thus 50 K, log g = 4.53 ± 0.14, and [Fe/H] = −0.03 ± 0.04. decided to attempt to measure the transit both in spectroscopy Following the calibration of Kervella et al. (2004), the V − K (Rossiter-McLaughlin effect) using the ELODIE spectrograph, colour implies a temperature of 4996 ± 40 K and the Strömgren and in photometry using the 1.20-m telescope on the same site. photometry gives 4950 ± 150 K, both estimates coherent The f/6 Newton focus of the 1.20-m telescope is equipped with with the spectroscopic temperature. Confronting the spectro- a CCD camera system (1024 × 1024 SITe back-illuminated scopic parameters obtained for HD 189733 with the Girardi CCD) giving a field of view of 11.8 arcmin size and a projected et al. (2002) stellar evolution models gives a radius of & pixel size of 0.69 arcsec. A filter wheel holds the filters U , B, R = 0.761 ± 0.014 R(, and a mass of M = 0.81 ± 0.03 M( using V, Rc and Ic. The observation sampling was limited by the read- the V magnitude, or M = 0.82 ± 0.025 M( using the K mag- out time of the CCD controller (90 s). Basic data processing and nitude. From the Baraffe et al. (1998) models, we find DAOPHOT stellar photometry (Stetson 1987) were applied to R = 0.75 ± 0.01 R( and M = 0.825 ± 0.025 M(. Kervella et al. the images. Aperture photometry of both brightest stars in the (2004) have calibrated a relation between the V − K colour field was performed, with an aperture of 14 pixels. Correction and radii measured from interferometry that shows very lit- of the sky background and cosmic removal were also applied. tle dispersion for low-mass stars (∼1%). This calibration gives The final light curves were then obtained using a single refer- a radius of R = 0.77 R( for HD 189733. Empirical and the- ence star, and applying an extinction correction fitted on the oretical estimates of the host star’s radius are thus in excel- data as a linear combination of the airmass. lent agreement, with a small error interval, thanks to the star’s position in a thin and slowly-evolving part of the lower main sequence and to the precision of the Hipparcos parallax. In the subsequent analysis we combine the above estimates into 3. Stellar characteristics of HD 189733 R = 0.76 ± 0.01 R( and M = 0.82 ± 0.03 R(. − HD 189733 (HIP 98505, GJ 4130) is a dwarf star in the north- A stellar rotation v sin i = 3.5 ± 1.0 km s 1 is estimated from ern hemisphere, listed in the Hipparcos catalog (ESA 1997) the calibration of the cross-correlation functions. Assuming with a visual magnitude V = 7.67, colour index B − V = 0.932, that the stellar spin axis is perpendicular to the line of sight, V−K = 2.13, and an astrometric parallax π = 51.94 ± 0.87 mas. we can derive the stellar rotational period of ∼11 days. The This puts the star at a close distance of 19.3 pc from the Sun chromospheric activity index S based on the relative flux level and allows us to derive a corresponding absolute magnitude of on CaII H and K lines was measured by Wright et al. (2004). MV = 6.25. Although the star is cataloged as G5, our analysis The value of S = 0.525 indicates a relatively active star. indicates a K1-K2 star. Following Santos et al. (2004), an LTE Table 1 lists the observed and derived parameters of the star high-resolution spectroscopic analysis of a high S/N spectrum HD 189733.

54 4.1. The discovery of a hot Jupiter transiting HD 189733

F. Bouchy et al.: The transiting hot Jupiter HD 189733 b. II. L17

Fig. 1. Radial-velocity measurements of HD 189733 superimposed Fig. 2. Phase-folded radial velocity measurements of HD 189733 su- on the best Keplerian solution. The higher density of points corre- perimposed on the best Keplerian solution. Error bars represent the sponds to spectroscopic-transit measurements. They are not used to photon-noise uncertainties. The inset shows a zoom near phase zero derive the Keplerian solution. Error bars represent the photon-noise where radial velocities exhibit the Rossiter-McLaughlin effect. uncertainties.

and planetary parameters. The somewhat large residuals around −1 4. Keplerian solution and spectroscopic transit the solution (15 ms ) are probably explained by the activity- induced jitter of the star. We checked (on the cross-correlation Radial velocity (RV) measurements of HD 189733 were con- functions) that the shape of the spectral lines was not varying ducted in August and September 2005 (from JD = 2 453 611 in phase with the radial-velocity change, as would be expected to 2 453 638). Figure 1 shows the RV measurements together in case of spot-induced RV variations. with the derived Keplerian solution. RV measurements made near phase zero are displayed in the inset of Fig. 2 and clearly 5. Photometric transit and characterisation shows the RV anomaly due to the Rossiter-McLaughlin effect of HD 189733 b (spectroscopic transit). A deviation from the Keplerian solution of about ±40 ms−1 occurs because the transiting planet occults From the ephemeris predicted by the radial velocities, 3 transit first the approaching limb and then the receding limb of the events (nights 2453629, 2453638 and 2453640) have been rotating star. The observation of this effect provides an unam- followed photometrically with the 1.20-m telescope at OHP. biguous confirmation of the transiting planet. HD 189733 is the The first night, a complete photometric transit was observed in third star known to present a Rossiter-McLaughlin effect due to the B band (Fig. 3). For the next 2 attempts, performed in the a planetary companion and actually the first to be identified as R band, only partial coverage of the transit was possible. We a transiting planet spectroscopically. Like HD 209458 (Queloz observed the transit egress on the first night and the ingress on et al. 2000a) and HD 149026 (Sato et al. 2005), HD 189733 the second. Because of non-optimum atmospheric conditions, presents a positive RV anomaly during the ingress phase of the the partial transit measurements are of poorer quality than the transit and a negative RV anomaly in the egress phase, indi- B-band observations, they are thus not used for our determi- cating that the stellar rotation is prograde relative to the planet nation of the planetary parameters. They are however of prime orbit. The symmetrical deviation seems to occur at mid-transit, importance to confirm the transit detection and precisely spec- indicating that the orbital plane is quite coplanar with the stellar ify the orbital period. equatorial plane. The amplitude of the anomaly is comparable The dispersion of the light curve in the B band (Fig. 3) is to the one measured on HD 209458 by Queloz et al. (2000a), in about 2 mmag at the beginning of the sequence and 3 mmag agreement with the fact that the star’s v sin i and the radius ra- at the end. The photon noise is 1.1 mmag, and the total dis- tio between the planet and its host star are very close for these persion is primarily due to the photon-noise on the compar- two systems. ison star and to the increasing airmass. A transit light curve The Keplerian solution is derived without the RV points ob- was fitted to the data using the mass and radius found for tained during the spectroscopic transit and using the constraint the host star in Sect. 3, the orbital parameters of Sect. 4 and of transit epochs given by the observed photometric transits. limb darkening coefficients in the B filter from Claret (2000) The best fit to the data yields a short-period orbit (P = 2.219 d) for Teff = 5000 K, log g = 4.5 and [M/H] = 0. The free param- with an eccentricity compatible with zero. The phase-folded eters are the transit central epoch, the radius ratio between radial-velocity curve is displayed in Fig. 2. The orbital ele- the star and planet, and the inclination angle of the orbit. We ments are listed in Table 1, jointly with the inferred stellar find Ttr = 2453629.3890± 0.0004, Rpl/R = 0.172 ± 0.003 and

55 Chapter 4. Main results

L18 F. Bouchy et al.: The transiting hot Jupiter HD 189733 b. II.

Fig. 4. Mass-radius diagram for the 9 transiting exoplanets. Jupiter and Saturn are indicated for comparison, as well as the loci of iso- densities at 0.3, 0.7 and 1.3 g cm−3. Data are from Pont et al. (2004) (OGLE-TR-111b), Konacki et al. (2005) (OGLE-TR-10b), Moutou et al. (2004) (OGLE-TR-132b), Torres et al. (2004) (OGLE-TR-56b), Bouchy et al. (2004) (OGLE-TR-113b), Laughlin et al. (2005) (HD 209458b and TrES-1), Sato et al. (2005) (HD 149026b). Fig. 3. Photometric transits of HD 189733 observed with the 1.20-m OHP telescope. Triangles correspond to the observation on B band made on September 15th 2005. Full circles and open circles corre- spond to observation on R band made respectively on September 24th and 26th 2005. The solid curve represents the best-fit model for the complete B-band transit.

i = 85.3 ± 0.1. The formal uncertainties are very small. The dominant source of error is likely to be the systematics in the photometry. To estimate their effect, we repeated the reduction of the photometry using different proceduresand different com- parison stars. This resulted in a change of 4% for the radius ra- tio and 0.3◦ for the inclination angle. The main parameters of the planet are therefore determined with remarkable accuracy, even from the “discovery” data, thanks to the very good deter- Fig. 5. Mass-period diagram for the 9 transiting exoplanets. mination of the primary star parameters. Orbital and physical parameters of HD 189733 b are listed in Table 1. Hipparcos (ESA 1997) observed HD189733 and obtained 176 reliable photometric measurements. Following Soderhjelm Although our programme is biased towards metal-rich (1999), Robichon & Arenou (2000) and Castellano et al. stars, the new candidate orbits a solar-metallicity star. Its short (2000), we looked for a transit signal in the Hipparcos data period tends then to weaken the proposed relation between sep- using the same shape and depth of the transit observed with aration and metallicity for hot Jupiters (Queloz et al. 2000b; the 1.20-m. Our best-fit period for the Hipparcos data is Sozzetti 2004). P = 2.218575± 0.000003 days. This result was also found and Figure 4 presents the mass-radius diagram of the 9 is described in details in Hébrard & Lecavelier des Etangs known transiting exoplanets. In terms of mass and radius, (2005). HD 189733 b is quite similar to the very hot Jupiters OGLE- TR-56b, 113b and 132b. 6. Summary and concluding remarks Figure 5 displays the period-mass diagram of the 9 known transiting exoplanets. HD 189733 b appears to be intermediate We have presented the characteristics of the new transiting hot between the class of very hot Jupiters and hot Jupiters and Jupiter in orbit around the star HD 189733, detected by the new provides in some way the missing link between planets from planet-search programme conducted with the ELODIE spectro- transit and radial-velocity surveys in terms of mass and period. graph. The period derived from the RV measurements is very HD 189733 b confirms the correlation between the periods (or short (P = 2.219 d) and the orbit is circular. The photometric orbital distances) and masses of transiting exoplanets pointed transit measurements allow the determination of the planetary out by Mazeh et al. (2005). At such a close distance from the mass (1.15 ± 0.04 MJ), radius (1.26 ± 0.03 RJ) and mean den- star, it is likely that HD 189733 b undergoes some evaporation sity (0.75 ± 0.08 g cm−3). (Lecavelier des Etangs et al. 2004; Baraffe et al. 2004).

56 4.1. The discovery of a hot Jupiter transiting HD 189733

F. Bouchy et al.: The transiting hot Jupiter HD 189733 b. II. L19

With the same visual magnitude as HD 209458 and even Charbonneau, D., Brown, T., Latham, D., et al. 2000, ApJ, 529, L45 brighter in infrared, HD 189733 belongs to the very short Claret, A. 2000, A&A, 363, 1081 list of bright stars with detected planetary transits. Only Da Silva, R., Udry, S., Bouchy, F., et al. 2005, A&A, in press HD209458b, TrES-1, HD149029b and HD189733b have ESA 1997, The HIPPARCOS and TYCHO catalogue, ESA-SP, 1200 parent stars brighter than V = 12. They therefore provide pri- Fischer, D. A., & Valenti, J. 2005, ApJ, 662, 1102 mary targets for additional ground-based and space-based mea- Girardi, L., Bertelli, G., Bressan, A., et al. 2002, A&A, 391, 195 Gonzalez, G. 1998, A&A, 334, 221 surements requiring very high signal-to-noise ratio observa- Hébrard, G., & Lecavelier des Etangs, A. 2005, A&A, submitted tions. Henry, G., Marcy, G., Butler, R., & Vogt, S. 2000, ApJ, 529, L41 Kervella, P., Thevenin, F., Di Folco, E., et al. 2004, A&A, 426, 297 Acknowledgements. We are grateful to all the night assistants and Konacki, M., Torres, G., Jha, S., et al. 2003, Nature, 421, 507 telescope staff of Observatoire de Haute Provence for their efforts Konacki, M., Torres, G., Sasselov, D., et al. 2005, ApJ, 624, 372 and their efficiency. We acknowledge X. Bonfils and F. Galland for Laughlin, G., Wolf, A., Vanmunster, T., et al. 2005, ApJ, 621, 1072 their help and the referee T. M. Brown for all his useful comments. Lecavelier des Etangs, A., Vidal-Madjar, A., McConnell, & J. C., We wish to thank the Programme National de Planetologie (PNP), Hébrard, G. 2004, A&A, 418, L1 the Swiss National Science Foundation (FNRS) and the Geneva Mazeh, T., Naef, D., Torres, G., et al. 2000, ApJ, 532, L55 University for their continuous support to our planet-search pro- Mazeh, T., Zucker, S., & Pont, F. 2005, MNRAS, 356, 955 grams. NCS would like to thank the support from Fundação para Moutou, C., Pont, F., Bouchy, F., et al. 2004, A&A, 424, L31 a Ciência e a Tecnologia (Portugal) the form of a scholarship (ref- Pont, F., Bouchy, F., Queloz, D., et al. 2004, A&A, 426, L15 erence SFRH/BPD/8116/2002) and a grant (reference POCI/CTE- Queloz, D., Eggenberger, A., Mayor, M., et al. 2000a, A&A, 359, L13 AST/56453/2004). R.D.S. would like to thank the support from Queloz, D., Mayor, M., Weber, L., et al. 2000b, A&A, 354, 99 Coordenação de Aperfeiçoamento de Pessoal de Nível Superior Robichon, N., & Arenou, F. 2000, A&A, 355, 295 (CAPES – Brazil) in the form of a scholarship. Sato, B., Fischer, D.A., Henry, G., et al. 2005, ApJ, in press Santos, N. C., Israelian, G., & Mayor, M. 2001, A&A, 373, 1019 Santos, N. C., Mayor, M., Naef, D., et al. 2002, A&A, 392, 215 References Santos, N. C., Israelian, G., & Mayor, M. 2004, A&A, 415, 1153 Santos, N. C., Israelian, G., Mayor, M., et al. 2005, A&A, 437, 1127 Alonso, R., Brown, T. M., Torres, G., et al. 2004, ApJ, 613, L153 Soderhjelm, S. 1999, Information Bulletin on Variable Stars, 4816, 1 Baraffe, I., Chabrier, G., Allard, F., et al. 1998, A&A, 337, 403 Sozzetti, A. 2004, MNRAS, 354, 1194 Baraffe, I., Selsis, F., Chabrier, G., et al. 2004, A&A, 419, L13 Stetson, P. B. 1987, PASP, 99, 191 Baranne, A., Queloz, D., Mayor, M., et al. 1996, A&AS, 119, 1 Torres, G., Konacki, M., Sasselov, D., 2004, ApJ, 609, 1071 Bouchy, F., Pont, F., Santos, N. C., et al. 2004, A&A, 421, L13 Udalski, A., Paczynski, B., Zebrun, K., et al. 2002, Acta Astron., 52, Bouchy, F., Pont, F., Melo, C., et al. 2005, A&A, 431, 1105 1 Castellano, T., Jenkins, J., Trilling, D. E., et al. 2000, ApJ, 532, L51 Wright, J. T., Marcy, G. W., Butler, R. P., et al. 2004, ApJS, 152, 261

57 Chapter 4. Main results

58 4.2. Other publications concerning the HD 189733 system

4.2 Other publications concerning the HD 189733 system

An important approach in the study of extrasolar planets is the direct observation of their infrared (IR) thermal emission, which can be done by observing secondary eclipses. The HD 189733 system offers a particular and very interesting situation for this kind of analysis: the relative small size of the star in comparison to the planet allows a maximum planet-star contrast (see Figure 8 of Fortney et al. 2006). Deming et al. (2006), by doing observations at 16 µm with the Spitzer telescope during a secondary eclipse, were able to determine the brightness temperature of the planet at this wavelength: T 16 = 1117 42 K, ± adopting the star-to-planet radius ratio from Bouchy et al. (2005).

In a series of photometric observations of the HD 189733 system using six telescopes around the world, Bakos et al. (2006b) refined the parameters of the planetary companion. A new determination of the planet’s radius (1.154 0.033 RJup) led to a higher density − ± ( 1000 kg m 3) and a higher brightness temperature at 16 µm(T 16 = 1279 90 K). ∼ ± Winn et al. (2006), by modelling the Rossiter-McLaughlin effect (see Figure 4.1), were able to measure the angle between the projected stellar spin axis and the projected plane- tary orbit normal for the HD 189733 system. They found λ = 1.4 1.1 ◦, which confirms so − ± our conclusion presented in Bouchy et al. (2005) concerning the coplanarity between the orbital plane and the stellar equatorial plane, suggesting that the inward migration process have preserved the spin-orbit alignment. Winn et al. (2007a) argued that the observation of the Rossiter-McLaughlin effect and the transit photometry, together with the measure- ment of the stellar rotation period, allow the determination of the true angle ψso between the stellar spin axis and the orbit normal, and not just the angle λso between their projec- ◦ tions. They determined an upper limit of 27 on the angle ψso for the HD 189733 system, which is the first extrasolar planetary system for which this angle has been measured.

The uncertainty in the planetary radius mostly depends on the determination of the stellar radius, since the size of the planet is measured relative to the size of its host star. Baines et al. (2007) have made a direct determination of the angular size of HD189733 (θ =0.377 0.024 mas) using interferometric observations, with no inferences of the stellar ± atmosphere. This leads to a physical stellar radius of 0.779 0.052 R!, which is slightly ± larger but compatible with previous determination.

Knutson et al. (2007b) observed the HD 189733 system with the Spitzer telescope, with observations that covered both a primary transit and the subsequent secondary transit, and were able to construct a map of the temperature distribution of the planet. They measured the day-night contrast of HD 189733 b, estimating a minimum and a maximum hemisphere- averaged brightness temperature at 8 µm, respectively 973 33 K and 1212 11 K, and ± ± a dayside brightness temperature of 1205.1 9.3 K. Their photometric data allowed a ± more precise determination of some the planet’s parameters than previous publications: a radius of 1.137 0.006 R , an orbital inclination of 85.61 0.04 ◦, and a planet-to-star ± Jup ± radius ratio of 0.1545 0.0002. ±

59 Chapter 4. Main results

Figure 4.1: Radial velocities of HD 189733 showing a model that fits both the Keplerian movement and the Rossiter-McLaughlin effect. Figure from Winn et al. (2006).

Using the InfraRed Spectrograph (IRS) on the Spitzer telescope, Grillmair et al. (2007) measured the spectrum of HD 189733 b at 7.5-14.7 µm during secondary eclipses, and found no indication of significant absorption by water or methane. Fortney & Marley (2007), however, propose the presence of water vapour absorption in the spectrum of HD 189733 b from 6.5 to 10 µm by examining the observations of Knutson et al. (2007b), which were conducted using the InfraRed Array Camera (IRAC), also on Spitzer. They suggested possible problems with the IRS spectra based on the expected presence of water absorption predicted by model atmospheres of extrasolar giant planets. Tinetti et al. (2007) analysed the transmission spectra of the HD 189733 system at three bands (3.6 µm, 5.8 µm, and 8 µm), and concluded that water vapour is indeed present in the atmosphere of this planet. Barnes et al. (2007) made near infrared observations (2.0-2.4 µm) of this system, and searched for H2O and CO, which are expected to have dominant opacities in this region. They found an upper limit that is not in agreement with their model, even if the model fits the mid-infrared Spitzer observations.

Pont et al. (2007b) obtained a very-high accuracy light curve from transit observations of the HD 189733 system using the HST. They made a direct geometric measurement of the orbital inclination angle (i = 85.68 0.04 ◦), the planet-to-star radius ratio (0.1572 ± ± 0.0004), and the ratio a/R =8.92 0.09, leading to improved values for the stellar and ! ± planetary radius, R =0.755 0.011 R! and R =1.154 0.017 R , respectively. They ! ± p ± Jup found no evidence for the presence of transiting satellite or second planet, planetary rings, planetary or stellar oblateness, neither transit timing variations. Their work also provided first insights into the geometry of spots on the stellar surface. At the same time, Moutou et al. (2007) conducted the first spectropolarimetric study of HD 189733, aiming to map the strength and topology of the large-scale magnetosphere of the star, and investigate the magnetic interaction between stars with close-in giant planets. They detected Zeeman signatures in all spectra (from 370 to 1000 nm), confirming the presence of magnetic fields at the stellar surface.

Croll et al. (2007) performed a search for other low mass companions in the HD 189733 system based on optical photometry observations using the MOST satellite. They used

60 4.2. Other publications concerning the HD 189733 system

Table 4.1: Summary of some published physical properties of the HD 189733 system, in 1 order of the publication year. All quoted references adopt a value of 0.82 0.03 M! for ± the stellar mass. I recalculated the density values assuming 1 R = 7.1492 107 m and Jup × 1M = 1.8986 1027 kg. Numbers in parentheses indicate the adopted reference for the Jup × respective parameter.

Rp R! Mp ρp i Tp Ref. −3 [RJup] [R"] [MJup] [kg m ] [deg] [K] 1 1.26 ± 0.03 0.76 ± 0.01 1.15 ± 0.04 710 85.3 ± 0.1 - 2 (1) (1) (1) 710 - 1117 ± 42b 3 1.154 ± 0.033 0.758 ± 0.016a (1) 930 85.79 ± 0.24 1279 ± 90b 4 1.10 ± 0.03 0.73 ± 0.02 1.13 ± 0.03 1050 86.1 ± 0.2 - 5 1.156 ± 0.046 0.753 ± 0.025 (4) 910 85.76 ± 0.29 - 6 1.19 ± 0.08 0.779 ± 0.052 (1) 850 - - 7 1.137 ± 0.006 0.757 ± 0.003 (1) 970 85.61 ± 0.04 1205.1 ± 9.3c 8 1.154 ± 0.017 0.755 ± 0.011 (1) 930 85.68 ± 0.04 -

1 Bouchy et al. (2005), 2 Deming et al. (2006), 3 Bakos et al. (2006b), 4 Winn et al. (2006), 5 Winn et al. (2007a), 6 Baines et al. (2007), 7 Knutson et al. (2007b), 8 Pont et al. (2007b) a value from Masana et al. (2006), b brightness temperature at 16 µm, c brightness temperature at 8 µm

the transit light curve to search for Earths, Super-Earths and Neptune sized planets in the interior regions of the HD 189733 b’s orbit. They ruled-out additional close-in exo- planets with sizes ranging from about 0.15 to 0.31 RJup. On the other hand, Bakos et al. (2006a) analysed the 2MASS1 data of HD 189733, and concluded that this system has a stellar companion separated by 11.2&& (216 AU), a M dwarf star with proper motion, radial velocity, and distance in common with the main star.

Recently, Henry & Winn (2008) have refined our estimate of the rotation period of HD 189733, obtaining a value of Prot = 11.953 0.009 days. In addition, they were able ◦ ± to set a lower limit of 54 on the inclination angle of the stellar rotation axis, which supports previous idea that the rotation axis of HD 189733 is not strongly misaligned with the normal to the orbital plane of HD 189733 b.

For purposes of comparison, Table 4.1 presents a summary of some physical properties of the HD 189733 system, showing the different estimates published by several authors. The discrepancies in the planet’s radius, which were determined based on observations collected in different epochs, may be the result of variability due to the relatively high level of the star’s magnetic activity. Figure 4.2 presents a mass-radius diagram with the 30 known transiting planets for which such parameters have been determined, comparing their respective mean densities with those of Jupiter, Saturn, Neptune, and Uranus.

1 Two Micron All Sky Survey (Cutri et al. 2003).

61 Chapter 4. Main results

Figure 4.2: Mass-radius diagram showing the 30 known transiting exoplanets, with Jupiter, Saturn, Neptune, and Uranus added for comparison. Dashed lines represent equidensities. Dotted lines on the right panel are isochrones from Baraffe et al. (1998, 2003) for ages of 0.5 Gyr (upper line) and 5 Gyr (lower line). The low mass stars plotted (open squares) are from S´egransan et al. (2003) and references therein.

4.3 Other hot Jupiters discovered with our programme

In addition to HD 189733 b, our programme has yielded the detection of another three hot Jupiters: HD 118203 b and HD 149143 b2 (Da Silva et al. 2006), and HD 185269 b2 (Moutou et al. 2006). They have periods between 4.1 and 6.8 days and minimum masses between 1.0 and 2.1 MJup. They are therefore expected to have a significant probability of presenting a transit event. I estimated the transit probability for this systems using Equa- tion 2.1, in which the stellar radii come from the luminosity-temperature-radius relation, luminosities and temperatures from Table A.1, and orbital elements from the respective papers of announcement (see Table 4.2). I also assumed Rp =0.1R!.

The paper of announcement of HD 118203 b and HD 149143 b (Da Silva et al. 2006) is presented in Subsection 4.3.1. HD 118203 b has a notably high transit probability of 18.9%, thanks to the large radius of the parent star (R! = 2.18 R!). However, the photometric measurements exhibited a constant signal at a 0.0047 mag level, and we found no evidence for a transit event. For HD 149143 b, the transit probability is 16.1%. Fischer et al. (2006) conducted photometric observations of this system, and they found a constant stellar flux at a 0.0003 mag level.

2 The planet around HD 149143 was independently discovered by Fischer et al. (2006) and the one orbiting HD 185269 was also published by Johnson et al. (2006).

62 4.3. Other hot Jupiters discovered with our programme

In Subsection 4.3.2 I present our paper of announcement concerning the hot Jupiter orbiting HD 185269 (Moutou et al. 2006). With a transit probability of 12.6%, HD 185269 b does not have a precise transit ephemeris due to its non-zero eccentricity (e =0.225), which induces uncertainties in the prediction of the central transit time. Nevertheless, Johnson et al. (2006) made a sequence of photometric observations of HD 185269, and concluded that this star is constant within 0.001 mag.

4.3.1 Paper: Hot Jupiters around HD 118203 and HD 149143

R. Da Silva, S. Udry, F. Bouchy, M. Mayor, C. Moutou, F. Pont, D. Queloz, N. C. Santos, D. S´egransan, and S. Zucker

Da Silva et al. 2006, Astronomy & Astrophysics 446, 717-722

1. Introduction

2. Sample and observational strategy

3. A short-period planet orbiting HD 118203

3.1. Stellar characteristics of HD 118203 3.2. Orbital solution for HD 118203 b 3.3. Low chromospheric activity for HD 118203

4. A hot Jupiter around HD 149143

4.1. HD 149143: stellar characteristics 4.2. Orbital solution for HD 149143 b

5. Summary and concluding remarks

63 Chapter 4. Main results

64 4.3. Other hot Jupiters discovered with our programme

A&A 446, 717–722 (2006) Astronomy DOI: 10.1051/0004-6361:20054116 & !c ESO 2006 Astrophysics

Elodie metallicity-biased search for transiting Hot Jupiters!

I. Two Hot Jupiters orbiting the slightly evolved stars HD 118203 and HD 149143

R. Da Silva1, S. Udry1, F. Bouchy2, M. Mayor1, C. Moutou2, F. Pont1, D. Queloz1, N. C. Santos3,1, D. Ségransan1, and S. Zucker1,4

1 Geneva Observatory, 1290 Sauverny, Switzerland e-mail: [email protected] 2 Observatoire de Marseille, France 3 Centro de Astronomia e Astrofísica da Universidade de Lisboa, Tapada da Ajuda, 1349-018 Lisboa, Portugal 4 The Weizmann Institute of Science, PO Box 26, Rehovot 76100, Israel

Received 29 August 2005 / Accepted 22 September 2005

ABSTRACT

We report the discovery of a new planet candidate orbiting the subgiant star HD 118203 with a period of P = 6.1335 days. The best Keplerian solution yields an eccentricity e = 0.31 and a minimum mass m2 sin i = 2.1 MJup for the planet. This star has been observed with the ELODIE fiber-fed spectrograph as one of the targets in our planet-search programme biased toward high-metallicity stars, on-going since March 2004 at the Haute-Provence Observatory. An analysis of the spectroscopic line profiles using line bisectors revealed no correlation between the radial velocities and the line-bisector orientations, indicating that the periodic radial-velocity signal is best explained by the presence of a planet-mass companion. A linear trend is observed in the residuals around the orbital solution that could be explained by the presence of a second companion in a longer-period orbit. We also present here our orbital solution for another slightly evolved star in our metal-rich sample, HD 149143, recently proposed to host a 4-d period Hot Jupiter by the N2K consortium. Our solution yields a period P = 4.09 days, a marginally significant eccentricity e = 0.08 and a planetary minimum mass of 1.36 MJup. We checked that the shape of the spectral lines does not vary for this star as well.

Key words. stars: individual: HD 118203 – stars: individual: HD 149143 – planetary systems – techniques: radial velocities

1. Introduction metallicity estimates. The programme mainly targets giant planets with short periods (Hot Jupiters). They are the ideal Stars hosting planets are significantly metal rich in compari- candidates in the search for photometric transits. From this son to field stars in the solar neighbourhood (Gonzalez 1997; survey, we present here a new short-period planet candi- Santos et al. 2001, 2003, 2005; Fischer & Valenti 2005). These date (P = 6.1335 d) with an eccentric orbit around the star authors have also shown that the probability of hosting a giant HD 118203. planet is a strongly rising function of the star metal content. According to their estimate, we can expect that up to 25−30% A similar planet-search programme was simultaneously started by the N2K consortium (Fischer et al. 2004) aiming of the more metal-rich stars ([Fe/H] > 0.2−0.3) host a giant planet. at the detection of short-period planets orbiting metal-rich stars. Fischer (2005; also in Fischer 2005) recently an- On the basis of this argument, a new survey was started nounced two new Hot-Jupiter detections around HD 149143 in March 2004 at the Haute-Provence Observatory with the and HD 109749. HD 149143 is amongst the stars already ob- high-precision ELODIE fiber-fed echelle spectrograph. The served in our sample and we present here our orbital descrip- main idea of this new programme is to bias our target sam- tion of the system. ple towards high-metallicity stars, much more likely to host planets. This will strongly increase our probability of finding Together, the N2K and ELODIE metallicity-biased planet- new planets in a sample of yet non-observed stars. The survey search programmes have detected five new Hot Jupiters in less uses the cross-correlation technique for the radial-velocity and than one year. One of them, HD 149026, is transiting in front of its parent star (Sato et al. 2005) allowing for the determination ! Based on radial velocities collected with the ELODIE spectro- of the planet radius and mean density. The planet is found to graph on the 193-cm telescope and photometric data on the 120-cm have an unexpectedly large core. This result clearly illustrates telescope, both at the Observatoire de Haute Provence, France. the importance of such programmes for our understanding of

65 Chapter 4. Main results

718 R. Da Silva et al.: Elodie metallicity-biased search for transiting Hot Jupiters. I.

planet interiors. However, when examining possible statistical Table 1. Observed and estimated parameters for HD 118203. (See text trends between orbital and stellar parameters to derive con- for references on the quoted values.) straints for planet formation models, we have to keep in mind the built-in bias of this subsample of exoplanets. In particular Spectral type K0 these planets must be removed when considering correlations V 8.05 with the star metallicity. B − V 0.699 The sample selection and observations of the new ELODIE π 11.29 ± 0.82 [mas] programme are described in Sect. 2. Stellar parameters, radial- M 3.31 velocity measurements of HD 118203 as well as the orbital so- V ± lution derived for the new Hot Jupiter candidate are presented Teff 5600 150 [K] in Sect. 3. This section also provides information on stellar ac- M! 1.23 ± 0.03 M( tivity and in particular the results of the bisector analysis for age 4.6 ± 0.8 Gyr the star. Section 4 reports our results for HD 149143 and con- log g –3.87 cgs clusions are presented in Sect. 5. [Fe/H] 0.10 ± 0.05 v sin i 4.7 [km s−1]

2. Sample and observational strategy

The star sample was selected from the HIPPARCOS catalogue 3. A short-period planet orbiting HD 118203 (ESA 1997). Initially, the bright northern dwarf stars (V < 8.5) of colours compatible with spectral types ranging from F8 3.1. Stellar characteristics of HD 118203 − to M 0 (0.45 < B V < 1.4) were considered. Then: HD 118203 (HIP 66192) is listed in the HIPPARCOS catalogue (ESA 1997) as a K0 dwarf in the northern hemisphere, with a i) evolved stars were removed from the list, using a criteria of visual magnitude V = 8.05, a colour index B − V = 0.699, and proximity to the main sequence in the HR diagram (2 mag); an astrometric parallax π = 11.29 ± 0.82 mas, setting a distance ii) all stars already members of the known main programmes of 88.6 pc from the Sun. The corresponding absolute magnitude of search for extra-solar planets in the northern hemisphere MV = 3.31 is too high for a K0 dwarf, indicating that the star were removed; is already slightly evolved and is in a subgiant stage. This is iii) we eliminated stars known as binaries from previous confirmed by the astrometric surface gravity log g = 3.87 dex CORAVEL measurements; estimated from the effective temperature and stellar mass (see iv) we also eliminated stars with high rotational velocities in below) and Hipparcos precise parallax (Santos et al. 2004). order to avoid active stars in the sample. In practice, stars From calibrations of the width and the surface of the with v sin i > 5 km s−1 were excluded from the sample. ELODIE cross-correlation functions (Santos et al. 2002; Naef 2003), we have estimated a projected rotation velocity The final selection comprises 1061 stars. They lie at distances v sin i = 4.7 km s−1 and a metallicity [Fe/H] = 0.10 for the star. up to 100 pc from the Sun. Comparing these observable stellar characteristics with the Observations were conducted with the high-precision stellar evolution models of Girardi et al. (2002), we can infer ELODIE fiber-fed echelle spectrograph (Baranne et al. 1996) the following intrinsic properties for HD 118203: a coarse esti- mounted on the 1.93-m telescope at the Haute-Provence mate of the effective temperature Teff = 5600 ± 150 K, a mass Observatory (France). The spectra have a resolution (λ/∆λ) of M! = 1.23 ± 0.03 M( and an age of 4.6 ± 0.8 Gyr. A similar about 42 000. Typical signal-to-noise ratios obtained in 20 min value of the effective temperature Teff = 5695 ± 50 K is also de- exposures range from ∼20 to 100 for the programme stars, cor- rived from the calibration in Santos et al. (2004) using the color responding to photon-noise errors between 5 and 20 ms−1 on index and metallicity of the star. The stellar parameters are individual measurements. The data reduction is performed on- gathered in Table 1. line during the observations by the automatic reduction soft- A check for stellar chromospheric activity was also per- ware (see Baranne et al. 1996, for a detailed description). formed by looking at the Ca  H absorption line in the spectra. After a single radial-velocity measurement, it is possible Figure 1 shows the co-added ELODIE spectra where the clear to obtain a very good estimate of the star metallicity by mea- absence of an emission feature indicates a low activity level, suring the surface of the cross correlation function (CCF) of as expected for slightly evolved subgiants (Wright 2004). The the ELODIE spectra (Santos et al. 2002; Naef 2003). The typi- activity-induced radial-velocity jitter is thus also expected to be cal uncertainty of the resulting metallicities is ∼0.05 dex com- low for this moderately rotating star. pared to values obtained from a high-resolution spectroscopic analysis. The expected percentage of >10−30% of giant plan- 3.2. Orbital solution for HD 118203 b ets orbiting stars with [Fe/H] ≥ 0.10 (Santos et al. 2001, 2004; Fischer & Valenti 2005) should lead to the discovery of a few ELODIE observations of HD118203 were conducted from tens of new planets in our programme, about ten of them in May 2004 (JD = 2 453 151) to July 2005 (JD = 2 453 553), short-period orbits. stimulated by the difference in radial velocity found among the

66 4.3. Other hot Jupiters discovered with our programme

R. Da Silva et al.: Elodie metallicity-biased search for transiting Hot Jupiters. I. 719

Table 2. ELODIE radial velocities of HD 118203. All data are rela- tive to the solar system barycentre. Given uncertainties correspond to photon-noise errors.

JD-2 400 000 RV Uncertainty [days] [km s−1] [km s−1] 53 151.4302 −29.263 0.013 53 153.4235 −29.231 0.014 53 155.3946 −29.640 0.011 53 217.3409 −29.523 0.011 53 218.3429 −29.320 0.011 53 219.3450 −29.243 0.013 53 220.3498 −29.215 0.012 53 224.3411 −29.300 0.009 53 226.3539 −29.152 0.017 53 227.3547 −29.232 0.026 53 228.3326 −29.533 0.012 53 229.3231 −29.549 0.011 Fig. 1. λ3968.5 Å Ca II H absorption line region of the summed 53 230.3545 −29.338 0.025 ELODIE spectra for HD 118203. No clear emission feature is ob- 53 231.3597 −29.257 0.013 served. For clarity, spectral features due to pollution by the thorium spectrum have been removed (e.g. around λ3966.5 Å). 53 232.3270 −29.209 0.018 53 233.3267 −29.261 0.009 53 234.3264 −29.474 0.010 first three observations. A set of 43 precise radial-velocity mea- 53 392.6777 −29.201 0.014 surements were then gathered. They are provided in Table 2. − A clear 6.1-d periodic variation is seen in this data set. 53 394.6422 29.598 0.027 The residuals around a Keplerian fit with this period are 53 396.6787 −29.257 0.015 however unexpectedly large (>40 ms−1) and shows a clear 53 397.6810 −29.153 0.015 additional radial-velocity drift as a function of the Julian 53 398.6756 −29.202 0.015 date. A simultaneous Keplerian + linear-drift adjustment 53 399.6799 −29.369 0.014 then provides a very satisfactory solution for the system 53 421.6764 −29.187 0.024 with a period P = 6.1335 ± 0.0006 days and an eccentricity 53 422.6439 −29.189 0.017 e = 0.309 ± 0.014. The slope of the linear drift is found to be − 49.7 ± 5.7 ms−1 yr−1. It is most probably accounted for by the 53 424.6543 29.556 0.013 presence of a second companion in the system, on a longer- 53 425.6554 −29.543 0.014 period orbit. This could explain the significantly non-zero value 53 426.6586 −29.351 0.012 found for the eccentricity at such a short period1. 53 428.6526 −29.190 0.025 A plot with the last-season radial-velocity measurements 53 429.5865 −29.166 0.028 of HD 118203 is shown in Fig. 2, together with the derived 53 430.5650 −29.446 0.016 solution. The residuals around the solution are displayed as 53 431.6340 −29.570 0.012 well in the bottom panel of the figure. The phase-folded radial- − velocity curve with the complete set of data points corrected for 53 486.5332 29.565 0.012 the observed linear drift is displayed in Fig. 3. The weighted 53 487.4809 −29.386 0.011 rms around the solution is σ(O−C) = 18.1 ms−1, slightly larger 53 488.4736 −29.260 0.012 than the individual photon-noise errors (∼15 ms−1). The or- 53 489.5168 −29.184 0.017 bital elements are listed in Table 3 with the inferred planetary 53 490.4763 −29.185 0.010 parameters. 53 491.4789 −29.298 0.011 53 492.4688 −29.594 0.011 3.3. Low chromospheric activity for HD 118203 53 517.4910 −29.547 0.011 Physical events in the stellar atmosphere, like the presence 53 519.4450 −29.212 0.011 of spots on the stellar surface, can change the observed 53 520.3896 −29.152 0.011 53 553.3845 −29.463 0.025 1 The period is however close to (higher than?) the observed cir- cularisation period determined from the sample of known exoplanets around 6 days (Halbwachs et al. 2004).

67 Chapter 4. Main results

720 R. Da Silva et al.: Elodie metallicity-biased search for transiting Hot Jupiters. I.

Fig. 2. Top: last 5-months radial-velocity measurements of HD 118203 Fig. 3. Top: phase-folded ELODIE radial-velocity measurements for superimposed on the best Keplerian + linear drift solution indicative the star HD 118203 after correction for the linear radial-velocity drift of a second longer-period companion in the system. Bottom: residuals observed in the data. Bottom: corresponding residuals around the so- around the solution. Error bars represent the photon-noise errors. lution. Error bars represent the photon-noise errors.

Table 3. ELODIE best Keplerian orbital solution obtained for HD 118203 as well as the inferred planetary parameters. For eccen- spectral-line profiles and induce a transient periodic radial- tric orbits T is defined as the time of the peri-astron passage. velocity signal similar to the one expected from the presence of a planet. The bisector analysis is one of the best tools to dis- P 6.1335 ±0.0006 [days] criminate between radial-velocity variations due to changes in ± − the spectral-line shapes and variations due to the real orbital T 53394.23 0.03 [JD 2 400 000] motion of the star (see Queloz et al. 2001, for a description of e 0.309 ±0.014 this method). V −29.387 ±0.006 [km s−1] The bisector inverse slope (BIS value) computed from the ω 155.7 ±2.4 [deg] HD 118203 spectra are plotted in Fig. 4 in comparison to the K 217 ±3 [m s−1]

corresponding radial-velocities. If the line-shape and the radial- Nmeas 43 velocity variations share the same origin, the BIS and radial σ(O−C) 18.1 [m s−1] velocities are expected to be correlated. This is not the case. Linear drift 49.7 ±5.7 [m s−1yr−1] Furthermore, no coherent signal is observed when the BIS are −5 phased with the orbital period of 6.1335 days. a1 sin i 11.61 [10 AU] −9 Moreover, data gathered during the still negative search f (m) 5.55 [10 M(] for a photometric transit (presented elsewhere) show that the m2 sin i 2.13 [MJup] star is constant at a 0.0047 mag level and thus further support a 0.07 [AU] its low-activity level. Consequently, the explanation of the ob- served radial-velocity variations due to stellar activity can be discarded. 4.1. HD 149143: stellar characteristics HD 149143 is listed in the Hipparcos catalogue (HIP 81022) = 4. A Hot Jupiter around HD 149143 with a G0 spectral type, a visual magnitude V 7.90, and a colour index B − V = 0.68. From calibrations of the width and During the preparation of this publication, another new Hot the surface of the ELODIE cross-correlation functions we have Jupiter candidate was identified in our programme, orbiting the estimated a projected rotation velocity v sin i = 3.9 km s−1 and star HD 149143. While gathering more points to characterize a metallicity [Fe/H] = 0.2. the system, we learned that the planet had just been announced The measured Hipparcos parallax of HD 149143 is by the N2K consortium (Fischer 2005; Fischer et al. 2005). π = 15.75 ± 1.07 mas setting the star 63.5 pc away from the We thus present here our ELODIE solution for the system, Sun. The inferred absolute magnitude is then MV = 3.88, also relying on the N2K photometric measurements for the transit slightly overluminous for a metallic G0 dwarf. This is con- non-detection. firmed by the astrometric estimate of log g = 4.10 dex. From the

68 4.3. Other hot Jupiters discovered with our programme

R. Da Silva et al.: Elodie metallicity-biased search for transiting Hot Jupiters. I. 721

Table 5. ELODIE radial velocities of HD 149143. All data are relative to the solar system barycentre.

JD-2 400 000 RV Uncertainty [days] [km s−1] [km s−1] 53 550.4467 12.226 0.013 53 554.4224 12.174 0.010 53 587.3709 12.232 0.014 53 589.3393 11.880 0.013 53 592.3736 12.102 0.012 53 611.3172 12.121 0.010 53 613.3130 11.992 0.010 53 614.3141 11.928 0.010

Table 6. ELODIE derived Keplerian orbital solution obtained for HD 149143 as well as the inferred planetary parameters. Both the quasi-circular and circular solutions are given because the derived ec- centricity is only marginally significant. For the eccentric orbit T is defined as the time of the peri-astron passage whereas for the circular orbit T indicates the maximum of radial velocities.

e free e = 0 P 4.088 ±0.006 4.089 ±0.006 [days] T 53588.0 ±0.4 53587.4 ±1.0 [JD − 2 400 000] e 0.08 ±0.04 0. V 12.059 ±0.004 12.056 ±0.003 [km s−1] Fig. 4. Radial velocities (RV, upper panel) and inverse bisector slope ω ± (BIS, middle panel) phased with the orbital period P = 6.1335 d for 42 35 0. [deg] − HD 118203. The independence of the two quantities is shown by the K 163 ±8 156 ±6 [m s 1] radial velocity vs. BIS plot (bottom panel). σ(O−C) 13.3 16.0 [m s−1] −5 a1 sin i 6.105 5.847 [10 AU] −9 Table 4. Observed and estimated parameters for HD 149143. (See text f (m) 1.82 1.59 [10 M(] for references on the quoted values.) m2 sin i 1.36 1.30 [MJup] a 0.052 0.052 [AU] Spectral type G0 V 7.9 B − V 0.68 4.2. Orbital solution for HD 149143 b π 15.75 ± 1.07 [mas] Although the number of ELODIE observations is small for MV 3.88 this star (Table 5), thanks to the large amplitude of the radial- −1 Teff 5730 ± 150 [K] velocity variation (K = 163 ms ), we can easily derive a quasi- circular Keplerian solution (e = 0.08 ± 0.04) with a period of M! 1.1 ± 0.1 M( 4.088 days, in agreement with the Fischer (2005) and Fischer age 7.6 ± 1.2 Gyr et al. (2005) announcement. The eccentricity of the orbit is log g 4.10 cgs small and marginally significant, we thus also provide the cor- [Fe/H] 0.20 ± 0.05 responding circular solution fixing e = 0 (Table 6). The derived v sin i 3.9 [km s−1] orbital elements coupled with the above estimate of the primary mass of 1.1 M( lead to a minimum mass m2 sin i = 1.36 MJup and a separation of 0.052 AU for the planetary companion. Figures 5 and 6 present the temporal and phased ELODIE models of Girardi et al. (2002) we obtain an effective temper- radial-velocity measurements superimposed on the derived ature Teff = 5730 ± 150 K, a primary mass M1 = 1.1 ± 0.1 M( Keplerian model as well as the residuals around the solu- and an age of 7.6 ± 1.2 Gyr. A calibration of the stellar ef- tion. The inferred orbital and planetary parameters are given fective temperature based on the stellar colour and metallicity in Table 6. (Santos et al. 2004, only slightly sensitive to the gravity of the As for the case of HD 118203, the bisector inverse slope star) leads then to Teff = 5790 ± 50 K, compatible with the value of the cross-correlation function has been calculated. No cor- given above. The stellar parameters are given in Table 4. relation is found between these bisector slopes and the radial

69 Chapter 4. Main results

722 R. Da Silva et al.: Elodie metallicity-biased search for transiting Hot Jupiters. I.

of 49.7 ms−1 yr−1 is observed, suggesting the presence of a second companion around the main star on a longer-period orbit. This additional companion could explain the somewhat high eccentricity of the orbit of this new planet at short period. We have also reported the ELODIE solution for another Hot Jupiter in our programme, HD149143, recently announced by the similar-goal N2K project. The derived best Keplerian quasi- circular solution presents a period of 4.09 days and the inferred planetary mass is 1.36 MJup. By selection these planets are orbiting metal-rich stars. They increase to five the number of Hot Jupiters detected in less than one year by dedicated metallicity-biased programmes. One of these candidates, HD 149026 (Sato et al. 2005), transits in front of its parent star and thus allows the determination of its radius and mean density when combining photometric and radial-velocity measurements. This demonstrates the efficiency of such approaches to find candidates suitable for constraining planet-interior models. However, the built-in biases of the sam- ple have to be kept in mind when examining possible statistical Fig. 5. ELODIE radial velocities of HD 149143 superimposed on the relations between the star metallicity and other orbital or stellar derived Keplerian quasi-circular model. Residuals around the solution parameters. are displayed in the bottom panel. Error bars represent the photon- noise errors. Acknowledgements. We thank the Swiss National Science Foundation (FNSRS) and the Geneva University for their continuous support to our planet-search programmes. We also thank the Haute-Provence Observatory for the granted telescope time. Support from Fundação para a Ciência e a Tecnologia (Portugal) to N.C.S. in the form of a scholarship (reference SFRH/BPD/8116/2002) and a grant (refer- ence POCI/CTE-AST/56453/2004) and support from Coordenação de Aperfeiçoamento de Pessoal de Nível Superior (CAPES - Brazil) to R.D.S. in the form of a scholarship are gratefully acknowledged as well.

References

Baranne, A., Queloz, D., Mayor, M., et al. 1996, A&AS, 119, 1 ESA 1997, The HIPPARCOS and THYCO catalogue, ESA-SP, 1200 Fischer, D. A., Valenti, J., & Marcy, G. 2004, in Stars as suns: activity, evolution and planets, Proc. IAU Symp., 219, 29 Fig. 6. Phase-folded ELODIE radial velocities of HD 149143 super- Fischer, D. A., & Valenti, J. 2005, ApJ, 622, 1102 imposed on the derived quasi-circular Keplerian solution. Error bars Fischer, D. A. 2005, in 10th anniversary of 51 Peg b Status of and represent the photon-noise errors. prospects for hot Jupiter studies, Conference at the Haute-Provence Observatory, August 2005 Fischer, D. A., et al. 2005, ApJ, in preparation velocities, excluding activity-induced variations of the shape of Halbwachs, J.-L., Mayor, M., & Udry, S. 2005, A&A, 431, 1129 the spectral lines as the source of the radial-velocity variations. Girardi, L., Bertelli, G., Bressan, A., et al. 2002, A&A, 391, 695 Gonzalez, G. 1997, MNRAS, 285, 403 Naef, D. 2003, Ph.D Thesis, Geneva University 5. Summary and concluding remarks Queloz, D., Henry, G. W., Sivan, J. P., et al. 2001, A&A, 379, 279 Santos, N. C., Israelian, G., & Mayor, M. 2001, A&A, 373, 1019 We have presented the characteristics of a new planet candidate Santos, N. C., Mayor, M., Naef, D., et al. 2002, A&A, 392, 215 in orbit around the subgiant star HD 118203, detected by the Santos, N. C., Israelian, G., Mayor, M., Rebolo, R., & Udry, S. 2003, new ELODIE planet-search programme biased towards metal- A&A, 398, 363 rich stars. The planet is in a rather eccentric orbit (e = 0.31), Santos, N. C., Israelian, G., & Mayor, M. 2004, A&A, 415, 1153 with a period of P = 6.1335 days, and is close to its parent star Santos, N. C., Israelian, G., & Mayor, M. 2005, A&A, 437, 1127 (a = 0.06 AU). Sato, B., Fischer, D., Henry, G. W., et al. 2005, ApJ, in press [arXiv:astro-ph/0507009] An additional trend of the radial-velocity measure- Wright, J. T. 2004, AJ, 128, 1273 ments increasing as a function of Julian date with a slope

70 4.3. Other hot Jupiters discovered with our programme

4.3.2 Paper: A hot Jupiter orbiting HD 185269

C. Moutou, B. Loeillet, F. Bouchy, R. Da Silva, M. Mayor, F. Pont, D. Queloz, N. C. Santos, D. S´egransan, S. Udry, and S. Zucker

Moutou et al. 2006, Astronomy & Astrophysics 458, 327-329

1. Introduction

2. Observations

3. Stellar characteristics of HD 185269

4. Planetary system

71 Chapter 4. Main results

72 4.3. Other hot Jupiters discovered with our programme

A&A 458, 327–329 (2006) Astronomy DOI: 10.1051/0004-6361:20066029 & !c ESO 2006 Astrophysics

ELODIE metallicity-biased search for transiting Hot Jupiters! III. A hot Jupiter orbiting the star HD 185269

C. Moutou1, B. Loeillet1, F. Bouchy2,3, R. Da Silva4, M. Mayor4, F. Pont4, D. Queloz4, N. C. Santos4,5,6, D. Ségransan4, S. Udry4, and S. Zucker7

1 Laboratoire d’Astrophysique de Marseille, Traverse du Siphon, 13013 Marseille, France e-mail: [email protected] 2 IAP, 98bis Bd Arago, 75014 Paris, France 3 Observatoire de Haute Provence, 04870 St Michel l’Observatoire, France 4 Observatoire de Genève, 51 ch. des Maillettes, 1290 Sauverny, Switzerland 5 Lisbon Observatory, Tapada da Ajuda, 1349-018 Lisboa, Portugal 6 Centro de Geofisica de Evora, Rua Romao Ramalho 59, 7002-554 Evora, Portugal 7 Weizmann Institute of Science, PO Box 26, Rehovot 76100, Israel

Received 13 July 2006 / Accepted 9 August 2006

ABSTRACT

We present new results of a search for extrasolar planets. The survey uses radial-velocity techniques and focuses on metal-rich stars. We used radial velocity measurements obtained with the echelle spectrograph ELODIE at the Observatoire de Haute Provence. New data have revealed a planetary companion to the slightly evolved star HD 185269. It belongs to the hot Jupiter class of planets, with a projected mass M sin i = 1.0 MJup and period P = 6.84 days. We describe the measurements that led to this discovery and discuss the orbital solution. Key words. stars: individual: HD 185269 – planetary systems – techniques: radial velocities

1. Introduction determinations of their mass and radius. However, non-biased target sample surveys are still needed in order to explore other There are approximately 40 known extrasolar planets with a parameters of the origins of planetary systems. semi-major axis smaller than 0.1 AU, the so-called hot Jupiters. In this paper, we report the detection of a new plane- Among them, seven have a mass of less than about 0.1 MJup and should better be called hot Neptunes. These hot Neptunes tary system around the main-sequence star HD 185269 at the could be very abundant as only recently have they become Observatoire de Haute Provence. In Sect. 2 the observations are available above the detection limit of intruments. 85% of these described; in Sect. 3 we derive the stellar parameters and the short-period planets have a metal-rich parent star, in comparison interpretation is discussed in Sect. 4. to 70% for the whole sample of extrasolar planets. There is thus an indication that the frequency of planets versus the metallicity of the parent star relationship is stronger for short-period planets (see also Sozzetti 2004). To understand whether this trend is sig- 2. Observations nificant, we need data from more short-period planets orbiting metal-rich stars. This is the goal of two observational programs Radial-velocity measurements of the main-sequence star using radial-velocity techniques; both surveys started in 2003 HD 185269 were obtained with ELODIE from June 2005 to and several planets have been discovered: HD 88133 (Fischer July 2006 in the framework of the survey of metal-rich stars as et al. 2005), HD 118203 (Da Silva et al. 2006), HD 149026 described in Da Silva et al. (2006). ELODIE is a fiber-fed echelle (Sato et al. 2005), HD 149143 and HD 109749 (Fischer et al. spectrograph (Baranne et al. 1996) mounted on the 193-cm 2006; Da Silva et al. 2006) and HD 189733 (Bouchy et al. telescope at the Observatoire de Haute Provence (France), de- 2005). The observational strategy of these programs, introduc- signed to obtain simultaneous spectra of a star and a calibration ing a bias in the selection of target stars based on their metallic- ThAr lamp, for precise radial-velocity measurements. ity, allows a more efficient detection of planets, and will clarify A first observation of HD 185269 shows its enrichment, the link between the metallicity of stars and the planet forma- with a metallicity deduced from the cross-correlation function tion and migration. Also, the discovery of short-period planets of 0.15. Further measurements of this star revealed fluctuations opens the possibility to discover new transiting systems (as e.g. in the radial velocity over a few days. The star was then regularly HD149026 and HD189733, see references above), with precise monitored. Individual measurements of HD185269 have errors of the order of 10−12 m s−1 (Table 1) while the standard devia- − ! Based on observations collected with the ELODIE spectrograph on tion of the series is about 90 m s 1. A 6.8-d periodicity is clearly the 1.93-cm telescope (Observatoire de Haute Provence, France). visible.

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328 C. Moutou et al.: New Hot Jupiter around HD 185269

Table 1. Radial-velocity measurements of HD 185269 obtained with Table 2. Stellar parameters of HD 185269. ELODIE from June 2005 to June 2006. Spectral Type G0 IV JD-2 400 000. VR error V 6.67 − − day km s 1 ms 1 B − V 0.58 53551.55350 0.535 0.011 Distance [pc] 47.37 53552.57440 0.639 0.013 MV 3.29 53553.54090 0.697 0.012 BC –0.06 53554.55910 0.690 0.008 Mbol 3.228 53555.54490 0.663 0.009 Teff [K] 5983 ± 62 53574.55610 0.687 0.012 log g 4.05 ± 0.22 53576.51360 0.631 0.011 [Fe/H] 0.10 ± 0.08 53577.53300 0.532 0.010 vsini [ km s−1] 5.5 53587.46980 0.672 0.011 Star mass [M(] 1.33 ± 0.07 53588.45620 0.677 0.009 Age[Gyr] 4.0 ± 1.0 53589.46500 0.606 0.012 53590.45090 0.631 0.014 53591.48810 0.530 0.011 53592.41560 0.543 0.011 53594.47180 0.672 0.010 53596.43110 0.677 0.010 53610.43430 0.652 0.010 53613.41040 0.593 0.008 53624.43320 0.630 0.008 53625.37140 0.551 0.010 53626.37360 0.508 0.008 53627.36100 0.597 0.008 53628.32940 0.672 0.008 53667.33630 0.502 0.008 53720.25140 0.617 0.010 53723.22240 0.603 0.011 53808.68270 0.640 0.012 53869.60890 0.672 0.008 53871.55960 0.538 0.008 Fig. 1. The CaII K absorption line of the ELODIE spectrum of 53872.61180 0.538 0.012 HD 185269, which shows no emission feature in its core. 53873.56570 0.589 0.015 53875.60370 0.716 0.010 53894.58770 0.648 0.010 53896.56560 0.683 0.009 53897.56800 0.681 0.009 53901.58760 0.661 0.014 53932.46560 0.609 0.015 53934.42560 0.518 0.009

3. Stellar characteristics of HD 185269 HD185269 has been observed by H (ESA 1997). It is a G0IV star with a visual magnitude V = 6.67 and color in- dex B − V = 0.58 and is located at 47.37 pc from the Sun. This results in an absolute V magnitude of 3.29. Strömgren narrow- band photometry on this star derives a metallicity of 0.02, i.e. al- most solar (Nordström et al. 2004). Comparatively, the metallic- ity derived from the cross-correlation function (0.15) is slightly higher. Direct spectroscopic analysis as described in Santos et al. (2004) allows us to derive: Teff = 5983 K, log g = 4.05 and [Fe/H] = 0.10 (Table 2). By comparison with the evolutionary tracks calculated by Girardi et al. (2002), we obtain a stellar mass of 1.33 ± 0.07 M( Fig. 2. The inverse slope of the bisector is shown against the radial- and an age of 4.0 ± 1.0 Gyr. The star is slightly evolved, as also velocity measurements of HD 185269. No correlation is observed. suggested by the lithium content in the spectrum. The projected rotational velocity derived from a calibration of the cross-correlation function is v sin i = 5.5 km s−1. The stellar chromospheric activity of HD 185269 is low, The bisector analysis of the cross-correlation functions as deduced from the absence of emission in the core of the shows no correlation with the measured velocimetric variations. CaII lines (Fig. 1). This excludes large amplitude radial-velocity The shape of stellar lines is therefore not related to the observed jitter due to intrinsic activity of the star. fluctuations in radial velocity (Fig. 2).

74 4.3. Other hot Jupiters discovered with our programme

C. Moutou et al.: New Hot Jupiter around HD 185269 329

Fig. 3. Radial-velocity measurements of HD 185269 superimposed on Fig. 4. Phase-folded radial velocity measurements of HD 185269 super- the best Keplerian solution. Error bars represent the photon-noise imposed on the best Keplerian solution. Error bars represent the photon- uncertainties. noise uncertainties. Bottom plot: residuals as a function of time.

Table 3. Parameters for the Keplerian solution and the planetary A short-period planet is a good target for transit detectabil- companion. ity; for HD 185269, the transit probability is about 8%. However, the eccentricity makes the error on the time reference large and Period [days] 6.8399 ± 0.0013 therefore the transit ephemeris is not precise. A photometric ob- Periastron epoch [JD-2453000] 797.152 ± 0.12 servation sequence will nevertheless be scheduled in the fol- Orbital eccentricity 0.225 ± 0.025 lowing months, in parallel with further radial-velocity measure- Radial velocity semi-amplitude [ m s−1] 93.56 ± 2.5 ments aimed at refining the orbital solution. Systemic velocity [ km s−1] 0.617 ± 0.002 O–C residuals [ m s−1] 16 Note added in proofs: After this paper was submitted, an inde- Orbital semi-major axis [AU] 0.077 pendent discovery of the same planet was announcedby Johnson Planet mass [MJ ] 1.03 ± 0.03 et al. (2006). Both sets of measurements and analysis are com- patible and no transit is visible.

4. Planetary system Acknowledgements. We thank the Programme National de Planetologie (PNP), the Swiss National Science Foundation (FNSRS) and the Geneva University for their continuous support of our planet-search programs. NCS thanks A careful check of stellar properties shows that the radial- the Fundação para a Ciência e a Tecnologia (Portugal) for the scholarship velocity fluctuations of HD 185269 were not likely of stellar ori- SFRH/BPD/8116/2002 and grant POCI/CTE-AST/56453/2004. gin. A Keplerian orbit was then adjusted on the 38 data points with photon noise uncertainty below 15 m s−1 (Fig. 3); the best parameters of the fit are a period of 6.84 days and an eccentricity References of 0.23. The final parameters of the orbital solution are presented Baranne, A., Queloz, D., Mayor, M., et al. 1996, A&AS, 119, 373 in Table 3. The residuals are 16 m s−1. Bouchy, F., Udry, S., Mayor, M., et al. 2005, A&A, 444, L15 With a stellar mass of 1.33 M(, the deduced minimum mass Da Silva, R., Udry, S., Bouchy, F., et al. 2006, A&A, 446, 717 is 1.03 ± 0.03 M for the planetary companion in an orbit with ESA 1997, VizieR Online Data Catalog, 1239, 0 Jup Fischer, D. A., Laughlin, G., Butler, P., et al. 2005, ApJ, 620, 481 a6.8399±0.0015 day period; the folded orbit is shown in Fig. 4. Fischer, D. A., Laughlin, G., Marcy, G. W., et al. 2006, ApJ, 637, 1094 Like HD 118203 (Da Silva et al. 2006), HD 185269 has a Girardi, L., Bertelli, G., Bressan, A., et al. 2002, A&A, 391, 195 short-period planet in eccentric orbit. The period of HD 185269 Halbwachs, J. L., Mayor, M., & Udry, S. 2005, A&A, 431, 1129 (6.84 days) is above the limit of circularization as expressed Johnson, J., Marcy, G., Fischer, D., et al. 2006, ApJ, in press Nordström, B., Mayor, M., Andersen, J., et al. 2004, A&A, 418, 989 in Halbwachs et al. (2005) and evidently, circular orbits are Santos, N. C., Israelian, G., & Mayor, M. 2004, A&A, 415, 1153 extremely rare when the period is larger than about 5 days Sato, B., Fischer, D. A., Henry, G. W., et al. 2005, ApJ, 633, 465 (Schneider 2006). Schneider, J. 2006, www.obspm.fr/planets

75 Chapter 4. Main results

76 4.4. Additional discoveries

4.4 Additional discoveries

Three planets with intermediate periods were also discovered with this programme. HD 43691 b and HD 132406 b (Da Silva et al. 2007) have periods P = 37 days and P = 974 days, respectively, and their parent stars were observed using both the ELODIE and SOPHIE spectrograph. The respective paper of announcement is presented in Subsec- tion 4.4.1. HD 45652 b (Santos et al. 2008, paper submitted) has period P = 43.7 days and was observed using ELODIE and CORALIE.

The subsample of 20 stars selected to be observed with the HARPS spectrograph has resulted, to date, in the detection of two low mass planetary companions, both having periods shorter than 5 days. HD 219828 b (Melo et al. 2007) is a hot Neptune with minimal mass Mp sin i = 19.8M⊕ and period P =3.8 days. The estimated transit probability is quite high: 18.3%. However, assuming for this planet the same radius of Neptune ( 0.35 M ), and using the stellar radius of 1.76 R! estimated by Melo et al. (2007), ∼ Jup the expected transit depth is only 0.04%, or 0.4 mmag. The other low mass planet is 3 HD 102195 b (Melo et al. 2007). The planet has a minimal mass Mp sin i =0.45 MJup and orbit its parent star with a period P =4.1 days. The transit probability is also significant (8.7%), but no transit was detected in the photometric observations of Ge et al. (2006). Our paper concerning this two planets is presented in Subsection 4.4.2.

Three stars selected in our original sample have planets discovered by other groups of search for extrasolar planetary systems. One of these planets has mass Mp sin i = 3 MJup and orbits the star HD 17156 with period P = 21 days and eccentricity e =0.67 (Fischer et al. 2007). The metallicity that we derived for this star using the calibration described in Subsection 3.3.2 is [Fe/H] = 0.04, which is smaller than the initial limit of 0.1 dex established by our observational strategy. For this reason, this star was not monitored with the ELODIE spectrograph. On the other hand, HD 17156 was subsequently monitored using SOPHIE observations, and the radial-velocity measurements confirmed the presence of the planet. The data are plotted in Figure 4.3, together with our best Keplerian solution fitted to the data. The orbital parameters from our solution agree with those determined by Fischer et al. (2007). Despite the relatively large orbital period, HD 17156 b has a favourable argument of periastron combined with a large eccentricity, which offers a high transit probability of 15.5%. Indeed, Barbieri et al. (2007) have recently made a transit detection when doing photometric observations of this system.

The other star from our main sample discovered by another group to harbour a plane- tary system is HD 155358 (Cochran et al. 2007). The radial velocities of this star indicate the presence of two planets: HD 155358 b has minimum mass Mp sin i =0.89 MJup, eccen- tricity 0.11, and period P = 195 days; HD 155358 c has minimum mass Mp sin i =0.5MJup, eccentricity 0.18, and period P = 530 days. A singular characteristic of this system is the very low value for the stellar metallicity, [Fe/H] = 0.68, which represents only 21% of the − 3 Firstly announced by Ge et al. (2006).

77 Chapter 4. Main results

Figure 4.3: Left: Radial-velocity measurements of HD 17156 gathered with the SOPHIE spectrograph together with our best Keplerian solution. The weighted rms around the solution is σ(O C) = 9.0 m s−1. Right: Phase-folded radial velocities and Keplerian − solution. Error bars represent the photon-noise errors. solar metal content. At the moment of the announcement of these planets, we had only one radial-velocity measurement of HD 155358 made using ELODIE. The reason why we not monitored this star is its low metallicity value (which we determined to be [Fe/H] = 0.78 based on our analysis of the gathered spectrum), remembering that we monitored − only metal-rich stars according to our observational strategy.

Robinson et al. (2007b) recently announced the discovery of a Jupiter-like planet in a Earth-like orbit around the star HD 75898. They derived for this planet a minimum mass Mp sin i =2.5MJup and a orbital period P =1.1 year. This star is also one of our target stars, and was observed using the ELODIE spectrograph. Figure 4.4 present our radial-velocity measurements together with the Keplerian solution derived by Robinson et al. (2007b). We gathered 21 measurements, but they were not enough to establish the one-year periodic variation. Even if they show a significant variation pic-to-pic (more than 160 m s−1), our measurements present a high dispersion (the weighted rms around the Keplerian solution is 23.5 m s−1), and the individual photon-noise errors are relatively large, with a mean value larger than 18 m s−1. All these factors, together with the poor coverage in phase, affected the interpretation of the data, making it difficult to identify the same periodic radial-velocity variation as that derived by Robinson et al. (2007b). No additional observations were made with the SOPHIE spectrograph.

Both HD 118203 b and HD 185269 b have short periods and eccentric orbits. It is worthwhile to notice that such non-zero eccentricities are consistent with their periods if we assume the limit P = 5 days of tidal circularisation derived by Halbwachs et al. (2005). The upper panel of Figure 4.5 shows the period-eccentricity diagram for the up-to-date list of known extrasolar planets4. We can notice that circular orbits are rare for periods larger than about 5 days. Planets with periods shorter than 5 days are normally expected to

4 Data from The Extrasolar Planets Encyclopaedia: http://exoplanet.eu

78 4.4. Additional discoveries

Figure 4.4: Radial-velocity measurements of HD 75898 collected with the ELODIE spec- trograph. The Keplerian orbital solution plotted here is the same fitted by Robinson et al. (2007b) to their data (Mpsin i= 2.5 MJup and P =1.1 yr). The bottom panel shows the residuals around the solution, and the weighted rms is σ(O C) = 23.5 m s−1. Error − bars represent the photon-noise errors. have their orbits circularised by tidal interactions with their parent star. However, this is not observed for some of the short-period planets in the period-eccentricity diagram, which may be explained by 1) perturbations by a possible distant companion, 2) the interaction with other planets in a multi-planetary system, and/or 3) the fact that the time for orbital circularisation has not yet been achieved.

In the lower panel of Figure 4.5 we can observe a well known property of the mass- period diagram: the lack of massive planets on short-period orbits. Some of the known extrasolar planets, however, have masses larger than 4 MJup and periods shorter than 100 days: some of them are possible members of binary systems; HD 162020 b, with a minimum mass of 13.8 MJup, is a possible brown dwarf (Udry et al. 2002); the most intriguing case is HAT-P-2 b (Bakos et al. 2007b), which is one of the transiting planets presented in Section 2.3. The mass-radius diagram of Figure 4.2 shows that this planet is in an intermediate position between the other known transiting planets and low-mass stars, and further studies are surely required for a better understanding of its nature.

For comparison, a summary of the main properties of the planets orbiting stars present in our sample, including those announced by other groups, are listed Table 4.2. The values for metallicity in this table are those from their respective publication when a spectroscopic analysis was done. Otherwise, the values are those from the calibration given by Equation 3.4.

79 Chapter 4. Main results

Figure 4.5: Upper panel: Period-eccentricity diagram for the list of known extrasolar planets. The dashed line indicates the maximum eccentricity of planets unaffected by tidal circularisation. Lower panel: Period-mass diagram for the same sample of stars. The dashed lines indicate the limits at 4 MJup and 100 days, and the dotted line connect the two massive objects around HD 168433. In both panels, filled circles represent planets in binary systems, and filled squares correspond to the planets orbiting stars present in our main sample. See text for more details on the quoted objects.

80 4.4. Additional discoveries

Table 4.2: Summary of the main orbital parameters of the systems discovered or confirmed with this work. Those announced by other groups are also shown.

P Mp sin i K a Planet name e −1 [Fe/H]! Ref. [days] [MJup] [m s ] [AU] HD 189733 b 2.2190 ± 0.0005 1.15 0.0 (fixed) 205 ± 6 0.0313 −0.03 1 HD 219828 b 3.8335 ± 0.0013 0.06 0.0 (fixed) 7.0 ± 2 0.048 0.19 9 HD 149143 b 4.088 ± 0.006 1.36 0.08 ± 0.04 163 ± 8 0.052 0.20 6, 3 HD 102195 b 4.1138 ± 0.0006 0.45 0.0 (fixed) 63 ± 0.5 0.05 0.05 5, 9 HD 118203 b 6.1335 ± 0.0006 2.13 0.309 ± 0.014 217 ± 3 0.07 0.15 3 HD 185269 b 6.8399 ± 0.0013 1.03 0.225 ± 0.025 93.6 ± 2.5 0.077 0.10 10, 8 HD 43691 b 36.96 ± 0.02 2.49 0.14 ± 0.02 124 ± 2 0.24 0.28 4 HD 45652 b 43.67 ± 0.08 0.60 0.25 ± 0.04 37.3 ± 1.6 0.24 0.15 12 HD 132406 b 974 ± 39 5.61 0.34 ± 0.09 115 ± 26 1.98 0.18 4

HD 17156 b 21.2 ± 0.3 3.12 0.67 ± 0.08 275 ± 15 0.15 0.24 7 HD 155358 b 195.0 ± 1.1 0.89 0.112 ± 0.037 34.6 ± 3.0 0.63 −0.68 2 HD 75898 b 418.2 ± 5.7 2.51 0.10 ± 0.05 58.2 ± 3.1 1.19 0.27 11 HD 155358 c 530.3 ± 27.2 0.5 0.176 ± 0.174 14.1 ± 1.6 1.22 −0.68 2

1 Bouchy et al. (2005), 2 Cochran et al. (2007), 3 Da Silva et al. (2006), 4 Da Silva et al. (2007), 5 Ge et al. (2006), 6 Fischer et al. (2006), 7 Fischer et al. (2007), 8 Johnson et al. (2006), 9 Melo et al. (2007), 10 Moutou et al. (2006), 11 Robinson et al. (2007b), 12 Santos et al. (2008)

81 Chapter 4. Main results

82 4.4. Additional discoveries

4.4.1 Paper: Two planets orbiting HD 43691 and HD 132406

R. Da Silva, S. Udry, F. Bouchy, C. Moutou, M. Mayor, J.-L. Beuzit, X. Bonfils, X. Delfosse, M. Desort, T. Forveille, F. Galland, G. H´ebrard, A.-M. Lagrange, B. Loeillet, C. Lovis, F. Pepe, C. Perrier, F. Pont, D. Queloz, N. C. Santos, D. S´egransan, J.-P. Sivan, A. Vidal-Madjar, and S. Zucker

Da Silva et al. 2006, Astronomy & Astrophysics 473, 323-328

1. Introduction

2. Observations

3. A planetary companion to HD 43691

3.1. Stellar characteristics of HD 43691 3.2. Orbital solution for HD 43691 b 3.3. Low chromospheric activity for HD 43691

4. A long-period planet orbiting HD 132406

4.1. Stellar characteristics of HD 132406 4.2. Orbital solution for HD 132406 b

5. Discussion and Conclusions

83 Chapter 4. Main results

84 4.4. Additional discoveries

A&A 473, 323–328 (2007) Astronomy DOI: 10.1051/0004-6361:20077314 & !c ESO 2007 Astrophysics

ELODIE metallicity-biased search for transiting Hot Jupiters! IV. Intermediate period planets orbiting the stars HD 43691 and HD 132406

R. Da Silva1, S. Udry1, F. Bouchy3, C. Moutou2, M. Mayor1, J.-L. Beuzit5, X. Bonfils5, X. Delfosse5, M. Desort5, T. Forveille5, F. Galland5, G. Hébrard3, A.-M. Lagrange5, B. Loeillet2, C. Lovis1, F. Pepe1, C. Perrier5, F. Pont1, D. Queloz1, N. C. Santos1,4, D. Ségransan1, J.-P. Sivan2, A. Vidal-Madjar3, and S. Zucker6

1 Observatoire Astronomique de l’Université de Genève, 1290 Sauverny, Switzerland e-mail: [email protected] 2 Laboratoire d’Astrophysique de Marseille, UMR6110 CNRS, Université de Provence, Traverse du Siphon, BP 8, 13376 Marseille Cedex 12, France 3 Institut d’Astrophysique de Paris, UMR7095 CNRS, Université Pierre & Marie Curie, 98bis bd Arago, 75014 Paris, France 4 Centro de Astrofísica, Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal 5 Laboratoire d’Astrophysique de Grenoble, BP 53X, 38041 Grenoble Cedex, France 6 Department of Geophysics and Planetary Sciences, Raymond and Beverly Sackler Faculty of Exact Sciences, Tel Aviv University, Tel Aviv 69978, Israel

Received 16 February 2007 / Accepted 4 July 2007

ABSTRACT

We report here the discovery of two planet candidates as a result of our planet-search programme biased in favour of high-metallicity stars, using the ELODIE spectrograph at the Observatoire de Haute Provence. One candidate has a minimum mass m2 sin i = 2.5 MJup and is orbiting the metal-rich star HD 43691 with period P = 40 days and eccentricity e = 0.14. The other planet has a minimum mass m2 sin i = 5.6 MJup and orbits the slightly metal-rich star HD 132406 with period P = 974 days and eccentricity e = 0.34. Additional observations for both stars were performed using the new SOPHIE spectrograph that replaces the ELODIE instrument, allowing an improved orbital solution for the systems. Key words. stars: individual: HD 43691 – stars: individual: HD 132406 – planetary systems – techniques: radial velocities

1. Introduction (Fischer et al. 2004), which monitors nearly 2000 main-sequence and subgiant stars. Another project was conducted by our team Following the first publications suggesting the metal-rich nature with the ELODIE spectrograph at the Observatoire de Haute of stars hosting giant planets in close orbits (Gonzalez 1997, Provence (Da Silva et al. 2006). From a sample of more than 1998), numerous studies on this theme have been published in a thousand solar-type stars, we selected the most metallic ones recent years (Santos et al. 2001, 2004; Fischer & Valenti 2005; after the first measurement to monitor their radial velocities. Gonzalez 2006). With the increasing number of known planets, Our programme has already yielded the detection of four Hot statistical studies were able to verify that the frequency of stars Jupiters, with periods between 2.2 and 6.8 days and minimum hosting a planetary companion is highly correlated with the stel- masses between 1.0 and 2.1 MJup, orbiting the stars HD 118203 lar metallicity. The results show that the probability of finding and HD1491431 (Da Silva et al. 2006), HD 189733 (Bouchy a close-in giant planet is about 25–30% for the most metal-rich et al. 2005), and HD1852691 (Moutou et al. 2006). / > stars ([Fe H] 0.3) and only 3% for stars with solar metallicity. For the star HD 189733, additional photometric measure- The extrasolar planet search biased in favour of high-metallicity ments have led to the observation of a planetary transit, making stars can thus rapidly lead to the discovery of planets in short- possible the determination of some parameters of the compan- < period orbits, the so-called Hot Jupiters (P 10 days), increas- ion, such as mass, radius and mean density (Bouchy et al. 2005; ing the chances of finding planetary transits. The identification Bakos et al. 2006; Winn et al. 2007). This system is our best re- of planets transiting bright stars provides a powerful approach to sult, a very good example of what we expect to obtain with our determine fundamental constraints on the mechanisms of planet programme, and revealed to be of particular interest for further formation, the physical properties of the exoplanet, and the ge- studies. Deming et al. (2006) analysed the infrared thermal emis- ometry of the system. sion during an eclipse of HD 189733 b using the Spitzer Space Based on these assumptions, a few programmes have Telescope (Werner et al. 2004) and determined the brightness been initiated with the aim of looking for planets orbiting temperature of the planet at 16 µm. Knutson et al. (2007), per- high-metallicity stars. One of them is the N2K consortium forming observations with Spitzer at 8 µm, were able to con- struct a map of the temperature distribution of HD 189733 b, ! Based on radial velocities collected with the ELODIE spectro- graph mounted on the 193-cm telescope at the Observatoire de Haute 1 The planet around HD 149143 was independently discovered by Provence, France. Additional observations were made using the new Fischer et al. (2006) and the one orbiting HD 185269 was also pub- SOPHIE spectrograph (run 06B.PNP.CONS) that replaces ELODIE. lished by Johnson et al. (2006).

85 Chapter 4. Main results

324 R. Da Silva et al.: ELODIE metallicity-biased search for transiting Hot Jupiters. IV.

estimating a minimum and a maximum brightness temperature Table 1. ELODIE and SOPHIE radial velocities of HD 43691. All val- at this wavelength. Fortney & Marley (2007) analysed the mid ues are relative to the solar system barycentre. The uncertainties corre- infrared observations of HD 189733 and suggested a possible spond to the photon-noise errors. presence of water vapor in the atmosphere of the planetary com- panion. Observations with the Hubble Space Telescope have also JD − 2 400 000 RV Uncertainty been proposed in order to perform precise measurements of the [days] [km s−1] [km s−1] size and the orbital inclination angle of HD 189733 b (Pont et al., ELODIE measurements in preparation). 53 333.6255 −29.123 0.011 In this paper, we report the discovery of two new planet can- 53 337.6057 −29.098 0.012 didates resulting from our ELODIE planet search programme bi- 53 398.4118 −29.015 0.012 ased towards metal-rich stars: a 2.5 Jupiter-mass planet orbiting 53 690.6694 −28.945 0.013 the star HD 43691 with period P = 40 days, and a 5.6 Jupiter- 53 692.6430 −28.962 0.013 mass planet in a long-period orbit of P = 974 days around the 53 693.6240 −28.983 0.014 star HD 132406. Such results are complemented by additional 53 714.5653 −28.912 0.014 measurements made using SOPHIE (Bouchy et al. 2006), the 53 715.5590 −28.903 0.015 new spectrograph that replaces ELODIE. 53 718.5549 −28.904 0.013 The radial velocity observations that have led to these results 53 719.5681 −28.885 0.010 − are described in Sect. 2. The observed and derived parameters 53 720.5323 28.853 0.017 53 721.5123 −28.892 0.018 of the star HD 43691 together with the orbital solution adopted 53 722.5237 −28.907 0.015 are presented in Sect. 3. The same are presented in Sect. 4 for 53 728.4315 −28.969 0.018 the star HD 132406. In Sect. 5 we discuss the present and future 53 749.5333 −28.989 0.014 status of the observational programme. 53 750.4982 −28.958 0.014 53 756.4444 −28.840 0.018 53 808.2808 −29.091 0.012 2. Observations 53 809.2850 −29.087 0.009 The HD 43691 and HD 132406 stars are both targets in our 53 839.3031 −28.960 0.019 − “ELODIE metallicity-biased search for transiting Hot Jupiters” 53 870.3405 28.881 0.017 53 872.3456 −28.923 0.020 survey (Da Silva et al. 2006), conducted from March 2004 un- til August 2006 with the ELODIE spectrograph (Baranne et al. SOPHIE measurements 1996) on the 193-cm telescope at the Observatoire de Haute 54 044.6274 −28.980 0.003 Provence (France). In this programme we essentially searched 54 051.6411 −28.830 0.004 for Jupiter-like planets orbiting metal-rich stars, assuming that 54 053.5968 −28.854 0.003 such stars are more likely to host giant planets. 54 078.6009 −29.059 0.004 − After obtaining the first spectrum of HD 43691 and 54 079.4859 29.040 0.004 54 080.4572 −29.018 0.004 HD 132406, we verified the high metallicity of these stars from 54 081.4440 −28.993 0.003 a calibration of the surface of the ELODIE cross-correlation 54 087.4744 −28.858 0.004 functions (Santos et al. 2002; Naef 2003). After three measure- 54 088.5858 −28.872 0.004 ments, we could clearly see in both stars a significant radial ve- 54 089.6137 −28.846 0.004 locity variation. We therefore conducted follow-up observations 54 142.4798 −29.064 0.004 with ELODIE, and we obtained 22 spectra of HD 43691 from 54 148.4607 −29.089 0.004 November 2004 (JD = 2 453 333) to May 2006 (JD = 2 453 872), 54 151.4010 −29.065 0.004 and 17 spectra of HD132406 from May 2004 (JD = 2 453 152) 54 155.4284 −28.981 0.004 to June 2006 (JD = 2 453 900). The ELODIE instrument was decommissioned in August 2006 and replaced by the SOPHIE spectrograph. Additional of astrometric signatures. Furthermore, the mass upper limits measurements were then obtained using this new instrument: produced by the IAD are in the stellar regime and therefore do 14 spectra of HD43691 from November 2006 (JD = 2 454 044) not provide any useful constraint. to February 2007 (JD = 2454155), and 4 spectra of In order to derive some of the fundamental stellar parame- HD132406 from December 2006 (JD = 2 454 080) to May 2007 ters, like effective temperature, surface gravity and metallicity, (JD = 2 454 230). using accurate spectroscopic analysis, we obtained a high S/N With ELODIE, the average signal-to-noise ratio (S/N) cal- spectrum (∼130 at λ5500 Å) of HD 43691 with the SOPHIE culated from the spectra at λ5500 Å is ∼40 for both stars, with spectrograph. a typical exposure time of 20 min. On the other hand, the gain in efficiency of SOPHIE compared to ELODIE is more than one order of magnitude in the high-resolution mode (used for high 3. A planetary companion to HD 43691 / precision radial-velocity measurements). Typical S N obtained 3.1. Stellar characteristics of HD43691 with SOPHIE are thus 2 times larger than those of ELODIE for exposure times 2–3 times smaller. Table 1 lists the radial veloci- HD 43691 (HIP 30057) is listed in the Hipparcos catalogue (ESA ties of HD 43691 and Table 2 lists those of HD 132406. 1997) as a G0 star in the northern hemisphere with visual mag- Following Zucker & Mazeh (2001), we tried to look for the nitude V = 8.03, color index B − V = 0.596, and parallax astrometric signatures of the two orbits in Hipparcos intermedi- π = 10.73 ± 1.16 mas (a distance of 93 pc from the Sun). ate astrometric data (IAD). HD132406, whose best-fit Keplerian The bolometric correction is BC = −0.034, derived from Flower period was close to the Hipparcos mission duration, seemed es- (1996). Using the Hipparcos parameters we derived an absolute pecially suitable for this kind of analysis. We found no evidence magnitude MV = 3.18, which represents a high luminosity for a

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Table 2. ELODIE and SOPHIE radial velocities of HD 132406. All val- Table 3. Observed and estimated parameters of HD 43691 and ues are relative to the solar system barycentre. The uncertainties corre- HD 132406. Some of the stellar parameters of HD 43691 were obtained spond to the photon-noise errors. from spectroscopic analysis while those of HD 132406 come from cal- ibrations of the ELODIE CCF. JD − 2 400 000 RV Uncertainty [days] [km s−1] [km s−1] HD 43691 HD 132406 ELODIE measurements Spectral type G0 IV G0 V 53 152.4773 −37.821 0.010 V 8.03 8.45 53 154.4825 −37.837 0.008 B − V 0.596 0.65 53 218.3605 −37.858 0.010 π 10.73 ± 1.16 14.09 ± 0.77 [mas] 53 520.4238 −37.928 0.013 MV 3.18 4.19 53 536.4140 −37.875 0.010 BC −0.034 −0.062 53 576.3681 −37.859 0.012 Teff 6200 ± 40 5885 ± 50 [K] 53 596.3805 −37.818 0.014 M! 1.38 ± 0.05 1.09 ± 0.05 M( 53 807.6666 −37.727 0.025 age 2.0–3.6 6.4 ± 0.8 Gyr 53 808.6537 −37.755 0.011 log g 4.28 ± 0.13 53 809.6643 −37.742 0.011 [Fe/H] 0.28 ± 0.05 0.18 ± 0.05 −1 53 869.5117 −37.779 0.011 v sin i 4.7 1.7 [km s ] 53 870.4406 −37.771 0.007 53 873.4386 −37.768 0.009 53 895.4251 −37.757 0.009 53 896.4387 −37.745 0.012 − 53 899.4286 37.743 0.019 . × −5 53 900.4376 −37.770 0.013 a false alarm probability of 1 3 10 using the approach de- scribed in Horne & Baliunas (1986). SOPHIE measurements In the top panel of Fig. 1 we plot the radial velocities of − 54 080.7252 37.718 0.003 HD 43691 and the Keplerian fit adopted using the two sets of 54 173.6848 −37.752 0.004 54 187.6332 −37.770 0.004 measurements. The middle panel shows the residuals around − the solution. The weighted rms around the solution is σE = 54 230.5803 37.790 0.004 −1 −1 17.5 m s for ELODIE, σS = 9.0 m s for SOPHIE, and −1 σES = 10.0 m s for the whole dataset. The bottom panel shows the phase-folded radial velocities. Table 4 lists the adopted or- G0 star. This suggests that HD 43691 is slightly evolved towards bital elements, together with the inferred planetary parameters. the subgiant branch. Nordström et al. (2004) found a difference of 1.19 mag from the ZAMS, indicating the degree of evolution of this star. 3.3. Low chromospheric activity for HD43691 Applying the spectroscopic analysis described in Santos et al. (2004) to the high S/N spectrum of HD43691 we ob- The radial velocity variations observed for a star can also be the result of physical events in the stellar atmosphere rather than tained: Teff = 6200 ± 40 K, log g = 4.28 ± 0.13, and [Fe/H] = 0.28 ± 0.05. From the calibrations of the ELODIE CCF the presence of an orbital companion. For example, spots on the (Santos et al. 2002; Naef 2003), we estimated a slightly smaller surface of an active star can change the observed spectral-line but comparable value for the metallicity ([Fe/H] = 0.22 ± 0.05), profiles and induce periodic variations in the measured radial and a projected rotation velocity vsin i = 4.7 km s−1. velocities. By analysing the line-bisector orientations one can With these stellar parameters, we estimated the mass and distinguish which of these situations is the real origin of the vari- age of HD 43691 using the Geneva models of stellar evolu- ations (Queloz et al. 2001). tion computed by Schaerer et al. (1993). We found a mass of The analysis of the line-bisector orientations, or bisector in- M! = 1.38±0.05 M( and an age between 2.0 and 3.6 Gyr, which verse slope (BIS) value, of HD 43691 shows that there is no cor- are in agreement with the determinations done by Nordström relation between the BIS values and the derived radial velocities et al. (2004): mass M! = 1.38 ± 0.08 M( and age 2.6 ± 0.5 Gyr. (Fig. 2, top panel). Thus the observed variations in radial veloc- These values are compatible with the star being slightly evolved, ity are not induced by stellar activity and rotation (as is the case especially taking into account the stellar metallicity (Mowlavi for HD 166435 in Queloz et al. 2001). The observed behaviour et al. 1998). The observed and derived stellar parameters of of the line bisectors also indicates that the radial velocity varia- HD 43691 are shown in Table 3. tions do not result from contamination by the light of a late-type binary companion (see e.g. the case of HD 41004 in Santos et al. 2002). Such variations are thus most probably due to the pres- 3.2. Orbital solution for HD43691b ence of a planetary companion orbiting HD 43691. The best Keplerian orbital solution fitted to the radial veloci- The chromospheric activity level can also be verified by ties of HD 43691, using both ELODIE and SOPHIE observa- means of the reemission in the core of Ca  absorption lines (e.g. tions, provides an orbit with period P = 36.96 ± 0.02 days λ3968.5 Å). By observing the respective spectral region in the and eccentricity e = 0.14 ± 0.02. With the estimated value for high S/N spectrum of HD 43691 obtained with SOPHIE (Fig. 2, the primary mass of 1.38 M(, we obtained a minimum mass bottom panel), we can note that this star is not active. Since this m2sin i = 2.49 MJup and a separation of 0.24 AU for the plan- line is located in the blue part of the spectral domain, where the etary companion. The solution includes a velocity zero-point of flux is appreciably lower, this kind of inspection requires a high the two datasets as a free parameter, and the difference between S/N spectrum, which is much better achieved by the higher effi- −1 them is ∆S−E = 23 ± 4 m s . For this solution, we estimated ciency of SOPHIE.

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326 R. Da Silva et al.: ELODIE metallicity-biased search for transiting Hot Jupiters. IV.

Table 4. Orbital elements for the best Keplerian solution of HD 43691 and HD 132406, as well as the inferred planetary parameters.

HD 43691 HD 132406 P 36.96 ± 0.02 974 ± 39 [days] T 54046.6 ± 0.5 53474 ± 44 [JD − 2 400 000] e 0.14 ± 0.02 0.34 ± 0.09 V −29.000 ± 0.003 −37.840 ± 0.008 [km s−1] ω 290 ± 5 214 ± 19 [deg] K 124 ± 2 115 ± 26 [m s−1] Nmeas 22 (E) + 14 (S) 17 (E) + 4 (S) −1 ∆S−E 23 ± 4 93 ± 17 [m s ] −1 σE 17.5 12.1 [m s ] −1 σS 9.0 4.1 [m s ] −1 σES 10.0 7.5 [m s ]

−4 a1sin i 4.17 9.73 [10 AU] −9 f (m) 7.06 1.30 [10 M(] m2sin i 2.49 5.61 [MJup] a 0.24 1.98 [AU]

4.2. Orbital solution for HD132406b A Keplerian solution fitted to the radial velocity measurements of HD 132406, from both ELODIE and SOPHIE observations, results in an orbit with period P = 974 ± 39 days and ec- centricity e = 0.34 ± 0.09. The velocity zero-point of the two datasets is a free parameter, and the difference between them is −1 ∆S−E = 93 ± 17 m s . The top panel of Fig. 3 shows the radial velocities of this star together with the best Keplerian solution. The bottom panel of the same figure shows the residuals around the adopted solution, −1 for which the weighted rms is σE = 12.1 m s for ELODIE, −1 −1 σS = 4.1 m s for SOPHIE, and σES = 7.5 m s for the com- bined set of measurements. HD 132406 is slightly fainter than Fig. 1. Top: ELODIE and SOPHIE radial velocities of HD 43691 plot- HD 43691, but has smaller photon-noise errors, which is prob- ted together with the best Keplerian solution that fits the combined mea- ably due to broader line profiles of HD 43691. Table 4 lists the −1 surements. Middle: Residuals around the solution, with σE = 17.5 m s orbital elements and the planetary parameters of the HD 132406 −1 −1 for ELODIE, σS = 9.0 m s for SOPHIE, and σES = 10.0 m s for the system, which was obtained with the best Keplerian fit. combined data points. Bottom: Phase-folded radial velocities with the As in the case of the star HD 43691, no correlation between best Keplerian solution. Error bars represent the photon-noise errors. the BIS values and the observed radial velocities is found for HD 132406 (Fig. 4). In addition, no chromospheric reemission is observed in the core of the Ca  absorption line at λ3968.5 Å.

5. Discussion and conclusions 4. A long-period planet orbiting HD 132406 In this paper we have announced the discovery of two new planet candidates resulting from our ELODIE search programme bi- 4.1. Stellar characteristics of HD132406 ased towards metal-rich stars. In this programme, a total of six planets have been discovered thus far, out of which four are Hot HD 132406 (HIP 73146) is listed in the Hipparcos catalogue as Jupiters (P < 10 days) and two are the intermediate-period plan- = . − = a G0 star with visual magnitude V 8 45, color index B V ets presented here. Five of the host stars have metallicity greater . π = . ± . 0 65, and parallax 14 09 0 77 mas (71 pc distant from the than 0.1 dex, while the one with a transiting very Hot Jupiter = . Sun). These parameters set a value of MV 4 19 for the absolute (HD 189733) is a solar-metallicity star. = − magnitude. The bolometric correction is BC 0.062. So far, we have observed almost 82% of the 1061 sample The metallicity and the projected rotation velocity of this star stars at least once, and for each one we estimated a value for are [Fe/H] = 0.18 ± 0.05 and vsin i = 1.7 km s−1 respectively, metallicity. We have verified that about 26% of the observed estimated from the calibrations of the ELODIE cross-correlation stars have [Fe/H] ≥ 0.1 dex (about 15% for 0.1 ≤ [Fe/H] < 0.2, functions. The effective temperature derived is Teff = 5885 ± 8% for 0.2 ≤ [Fe/H] < 0.3, and 3% for 0.3 ≤ [Fe/H] < 0.4). 50 K and comes from the calibrations of Teff as a function of Furthermore, according to the percentage of stars with planets B − V and [Fe/H] from Santos et al. (2004). The derived mass per metallicity bin determined by Santos et al. (2004), 10−30% and age are M! = 1.09 ± 0.05 M( and 6.4 ± 0.8 Gyr, from the of the stars with [Fe/H] ≥ 0.1 dex are likely to host a giant planet Geneva models of stellar evolution. These parameters are listed (about 9, 24 and 28% respectively for the same three ranges of in Table 3. metallicity mentioned above). Applying these percentages to the

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Fig. 3. Top: ELODIE and SOPHIE radial velocities of HD 132406 plot- ted together with the adopted Keplerian solution that better fits the combined measurements. Bottom: Residuals around the solution, with −1 −1 σE = 12.1 m s for ELODIE, σS = 4.1 m s for SOPHIE, and −1 σES = 7.5 m s for the combined data points. Error bars represent the photon-noise errors.

Fig. 2. Top: Comparison between bisector inverse slope (BIS) and radial velocities of HD 43691 showing no correlation between them for both ELODIE and SOPHIE measurements. Bottom: λ3968.5 Å Ca  absorp- tion line region of the high S/N spectrum obtained for HD 43691. No clear emission feature is observed in the centre of this line, indicating a low activity level.

Fig. 4. Comparison BIS and radial velocities of HD 132406 showing no correlation between them. Both ELODIE and SOPHIE measurements 867 observed stars, we conclude that the number of giant planets are plotted. we expect to discover in each of those metallicity ranges is 12, 17 and 7 respectively, a total of 36 planets, from which roughly 25% (9 planets) are predicted to be Hot Jupiters. The stars presented in this paper were also observed using Although most of our target stars were already observed, the new SOPHIE spectrograph. The proposed orbital solutions, only 75% of the metal-rich stars have a minimum of three mea- first found with ELODIE, were improved with the new observa- surements. Stars with one or two spectra need more observa- tions. With ELODIE decommissioned, this new instrument will tions before being rejected or classified as possible planet hosts. also continue monitoring other high-metallicity stars, especially On the other hand, long-period planets are more difficult to de- the most promising cases. tect, and stars showing long-term radial-velocity trends also need Acknowledgements. We thank the Swiss National Science Foundation (FNSRS) more observations. In any case, we have already found almost a and the Geneva University for their continued support to our planet-search pro- half of the expected number of Hot Jupiters among the metallic grammes, and the Observatoire de Haute Provence for the granted telescope portion of our sample. time. N.C.S. would like to thank the support from Fundação para a Ciência

89 Chapter 4. Main results

328 R. Da Silva et al.: ELODIE metallicity-biased search for transiting Hot Jupiters. IV.

e a Tecnologia (FCT), Portugal, in the form of a grant (reference POCI/CTE- Fischer, D. A., Laughlin, G., Marcy, G. W., et al. 2006, ApJ, 637, 1094 AST/56453/2004). This work was supported in part by the EC’s FP6 and by FCT Flower, P. J. 1996, ApJ, 469, 355 (with POCI2010 and FEDER funds), within the HELAS international collabo- Fortney, J. J., & Marley, M. S. 2007, ApJL, submitted ration. The support from Coordenação de Aperfeiçoamento de Pessoal de Nível Gonzalez, G. 1997, MNRAS, 285, 403 Superior (CAPES - Brazil) to R.D.S. in the form of a scholarship are gratefully Gonzalez, G. 1998, A&A, 334, 221 acknowledged as well. Gonzalez, G. 2006, PASP, 118, 1494 Horne, J. H., & Baliunas, S. L. 1986, ApJ, 302, 757 Johnson, J. A., Marcy, G. W., Fischer, D.A, et al. 2006, ApJ, 652, 1724 References Knutson, H. A., Charbonneau, D., Allen, L. E., et al. 2007, Nature, 447, 183 Moutou, C., Loeillet, B., Bouchy, F., et al. 2006, A&A, 458, 327 Baranne, A., Queloz, D., Mayor, M., et al. 1996, A&AS, 119, 373 Mowlavi, N., Meynet, G., Maeder, A., Schaerer, D., & Chatbonnel, C. 1998, Bakos, G. A., Knutson, H., Pont, F., et al. 2006, ApJ, 650, 1160 A&A, 335, 573 Bouchy, F., & the SOPHIE team 2006, SOPHIE: the Successor of the Naef, D. 2003, Ph.D. Thesis, Geneva University Spectrograph ELODIE for Extrasolar Planet Search and Characterization. Nordström, B., Mayor, M., Andersen, J., et al. 2004, A&A, 418, 989 In Tenth Anniversary of 51 Peg-b: Status of and Prospects for Hot Jupiter Queloz, D., Henry, G. W., Sivan, J. P., et al. 2001, A&A, 379, 279 Studies. Colloquium held at OHP, France, August 22–25, 2005. ed. L. Arnold, Santos, N. C., Israelian, G., & Mayor, M. 2000, A&A, 363, 228 F. Bouchy, & C. Moutou (Paris: Frontier Group), 319. Santos, N. C., Israelian, G., & Mayor, M. 2001, A&A, 373, 1019 Bouchy, F., Udry, S., Mayor, M., et al. 2005, A&A, 444, L15 Santos, N. C., Mayor, M., Naef, D., et al. 2002, A&A, 392, 215 Da Silva, R., Udry, S., Bouchy, F., et al. 2006, A&A, 446, 717 Santos, N. C., Israelian, G., & Mayor, M. 2004, A&A, 415, 1153 Deming, D., Harrington, J., Seager, S., & Richardson, L. J. 2006, ApJ, 644, 560 Schaerer, D., Charbonnel, C., Meynet, G., Maeder, A., & Schaller, G. 1993, ESA 1997, The Hipparcos and Tycho Catalogue, ESA SP-1200 A&AS, 102, 339 Fischer, D. A., & Valenti, J. 2005, ApJ, 622, 1102 Werner, M. W., Roellig, T. L., Low, F. J., et al. 2004, ApJS, 154, 1 Fischer, D. A., Valenti, J., & Marcy, G. 2004, in Stars as Suns: Activity, Winn, J. N., Holman, M. J., Henry, G. W., et al. 2007, ApJ, 133, 1828 Evolution, and Planets, Proc. IAU Symp., 219, 29 Zucker, S., & Mazeh, T. 2001, ApJ, 562, 549

90 4.4. Additional discoveries

4.4.2 Paper: Two hot companions around HD 102195 and HD 219828

C. Melo, N. C. Santos, W. Gieren, G. Pietrzynski, M. T. Ruiz, S. G. Sousa, F. Bouchy, C. Lovis, M. Mayor, F. Pepe, D. Queloz, R. Da Silva, and S. Udry,

Melo et al. 2007, Astronomy & Astrophysics 467, 721-727

1. Introduction

2. Observations

3. A new planet around HD 219828

3.1. Stellar characteristics 3.2. HARPS orbital solution

4. The planet around HD 102195

4.1. Stellar characteristics of HD 102195 4.2. The Harps orbital solution 4.3. Correcting the radial velocities for the stellar noise

5. Concluding remarks

91 Chapter 4. Main results

92 4.4. Additional discoveries

A&A 467, 721–727 (2007) Astronomy DOI: 10.1051/0004-6361:20066845 & !c ESO 2007 Astrophysics

A new Neptune-mass planet orbiting HD 219828!

C. Melo1, N. C. Santos2,3,4, W. Gieren5, G. Pietrzynski5, M. T. Ruiz6, S. G. Sousa2,7,8, F. Bouchy9, C. Lovis3, M. Mayor3, F. Pepe3, D. Queloz3, R. da Silva3, and S. Udry3

1 European Southern Observatory, Casilla 19001, Santiago 19, Chile e-mail: [email protected] 2 Centro de Astronomia e Astrofísica da Universidade de Lisboa, Observatório Astronómico de Lisboa, Tapada da Ajuda, 1349-018 Lisboa, Portugal 3 Observatoire de Genève, 51 ch. des Maillettes, 1290 Sauverny, Switzerland 4 Centro de Geofisica de Évora, Rua Romão Ramalho 59, 7002-554 Évora, Portugal 5 Universidad de Concepcion, Departamento de Fisica, Casilla 160-C, Concepcion, Chile 6 Departamento de Astronomia, Universidad de Chile, Casilla Postal 36D, Santiago, Chile 7 Centro de Astrofísica da Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal 8 Departamento de Matemática Aplicada, Faculdade de Ciências da Universidade do Porto, Portugal 9 Institut d’Astrophysique de Paris, 98bis Bd. Arago, 75014 Paris, France Received 30 November 2006 / Accepted 28 January 2007 ABSTRACT

Two years ago a new benchmark for the planetary survey was set with the discoveries of three extrasolar planets with masses be- low 20 M⊕. In particular, the serendipitous discovery of the 14 M⊕ planet around µ Ara found with HARPS with a semi-amplitude of only 4 m s−1 put in evidence the tremendous potential of HARPS for the search of this class of very low-mass planets. Aiming to discovering new worlds similar to µ Ara b, we carried out an intensive campaign with HARPS to observe a selected sample of northern stars covering a range of metallicity from about solar to twice solar. Two stars in our program were found to present radial velocity variations compatible with the presence of a planet-mass companion. The first of these, HD 219828, was found to be orbited by a planet with a minimum mass of 19.8 M⊕ and an orbital period of 3.83 days. It is the 11th Neptune-mass planet found so far orbiting a solar-type star. The radial velocity data clearly show the presence of an additional body to the system, likely of planetary mass. The second planet orbits HD 102195, has a mass of 0.45 MJup and an orbital period of 4.11 days. This planet has been already announced by Ge et al. (2006, ApJ, 648, 683). Our data confirm and improve the orbital solution found by these authors. We also show that the high residuals of the orbital solution are caused by stellar activity, and use the bisectors of the HARPS cross-correlation function to correct the noise introduced by stellar activity. An improved orbital solution is obtained after this correction. This kind of analysis may be used in the future to correct the radial-velocities for stellar activity induced noise. Key words. stars: individual: HD 219828 – stars: planetary systems – planetary systems: formation – techniques: radial velocities

1. Introduction of µ Ara was only made possible thanks to the two main fac- tors. First, the intrinsic stability of an instrument like HARPS The serendipitous discovery by Santos et al. (2004a) of a 14 M⊕ which is able to keep a long-term (several years) radial velocity planet orbiting µ Ara, simultaneously followed by the announce- accuracy better than 1 m s−1. Second, the observational strategy. ment of another planet of similar mass around 55 Cnc (McArthur µ Ara was observed in a context of a asteroseismology program, et al. 2004) and of a 21 M⊕ mass planet around GJ 436 b (Butler where a unique target was monitored during the whole night and et al. 2004) set a new benchmark for planet surveys. Since then, a a large number of spectra was collected with exposure times of few more cases have been announced(Rivera et al. 2005; Bonfils about a 2 min, in order to measure the stellar acoustic oscilla- et al. 2005; Vogt et al. 2005; Udry et al. 2006), including a sys- tion modes. Seen from the perspective of a radial velocity plan- tem of three Neptune-mass bodies (Lovis et al. 2006). For the etary search, these oscillations are unwanted and are considered first time in the literature, the minimum masses of a planet found as stellar noise. This intrinsic noise is well illustrated by the case by Doppler shift techniques were indicated in Earth-masses and of µ Ara: during the nights when the star was measured with not in Jupiter- or Saturn-masses. Even more exciting was the HARPS over a several hours period, the residuals around the or- fact that based on their orbital characteristics and masses, it is bital fit were merely around 0.4 m s−1, a value that increased to strongly suggested that these planets may be rocky or icy in na- near 1.5 m s−1 for the remaining nights, when the radial-velocity ture (Alibert et al. 2006). These discoveries may have thus un- of the star was an average of only about 15 min integrations veiled the first super-earths, and shown that telluric planets may (Santos et al. 2004a); this value would be around 3 m s−1 if be common in the solar neighborhood. only one single (short exposure) radial-velocity measurement The detection of these kinds of planets is not easy to be was done. In other words, to be able to achieve the 1 m s−1 pre- achieved. For example, the discovery of the 14 M⊕ companion cision one needs to average the radial velocities for a given star

! over periods of several minutes (at least 15−20 min), making the Based on observations collected at the La Silla Parana Observatory, searches for very low mass planets very time consuming (see the ESO (Chile) with the HARPS spectrograph at the 3.6 m telescope (Observing runs 075.C-0332, 076.C-0155, and 77.C-0101). seismological results on µ Ara by Bouchy et al. 2005).

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722 C. Melo et al.: A new Neptune-mass planet orbiting HD 219828

Surprisingly enough, µ Ara and 55 Cnc were back then the Table 1. Stellar parameters for HD 219828. only two stars that were (incidentally) measured with enough precision so that such low mass planets could have been de- Parameter Value Reference tected. This may imply that very low mass short period planets Spectral type G0IV Hipparcos are very common. This result is supported by recent simulations Parallax [mas] 12.33 ± 1.01 Hipparcos (Ida & Lin 2004; Benz et al. 2006) that suggest that very low- Distance [pc] 81.1 Hipparcos mass planets may be more frequent than the previously found mv 8.04 Hipparcos giant worlds. B − V 0.654 Hipparcos Mv 3.49 – These theoretical results and the discovery of several Luminosity [L ] 3.34† – Neptune-mass planets suggested that if solar type stars are cor- ( Mass [M(] 1.24 Schaerer et al. (1993) rectly monitored, the number of detection of such low-mass & log RHK −5.04 HARPS rocky planets can be very high. Having this in mind, we started v sin i [km s−1] 2.9†† – a project to use the HARPS spectrograph to monitor a sample Teff [K] 5891 ± 18 – of about 40 northern stars with metallicity ranging from solar to log g 4.18 ± 0.02 – about twice solar. In the present paper, we present the discov- ξt 1.18 ± 0.02 – ery of a new Neptune-mass planet orbiting the star HD 219828, [Fe/H] +0.19 ± 0.03 – as well as the confirmation of a Jovian-mass planet orbiting † Using the bolometric correction of Flower (1996). HD 102195. We also find evidence for the existence of a longer †† From HARPS spectra using a calibration similar to the one presented period planet companion to the former star. The presentation of by Santos et al. (2002). the whole sample, along with a detailed study of the behavior of the radial velocities with other properties (e.g. spectral type and activity level) is postponed to a future paper. mass is used in conjunction with the stellar parallax (e.g. Santos et al. 2004b). This result is compatible with the spectral clas- sification of the star (G0IV – ESA 1997), and indicates that 2. Observations HD 219828 may be slightly evolved. The stellar radius estimated The observations were carried out using the HARPS spectro- using the luminosity–temperature–radius relation is 1.76 R(. graph (3.6-m ESO telescope, La Silla, Chile), in three different From the HARPS spectra we derived both a chromospheric & = − . observing runs between May 2005 and July 2006 (ESO observ- activity index (log RHK 5 04) and an estimate of the projected −1 ing runs 075.C-0332,076.C-0155,and 77.C-0101).Further mea- rotational velocity of the star (v sin i = 2.9 km s ). From the ac- surements of the two stars discussed in this paper were also done tivity level and the B − V colour we derive an age of 6.5Gyr in different HARPS GTO runs, in collaboration with the Geneva (Henry et al. 1996) (at least above 2Gyr – Pace & Pasquini team. 2004), and a rotational period of 26 days (Noyes et al. 1984). Radial-velocities were obtained directly using the HARPS All these values suggest that HD 219828 is an old chromospher- pipeline. We refer to Pepe et al. (2004) for details on the data re- ically quiet star. According to the Hipparcos catalog, this star is duction. The individual spectra were also used to derive both the considered to be constant in photometry, with a scatter of only Bisector Inverse Slope (BIS) of the HARPS Cross-Correlation 0.013 mag (typical for a constant star of its magnitude). Function (CCF), as defined by Queloz et al. (2001), as well & as a measurement of the chromospheric activity index log RHK, 3.2. HARPS orbital solution following the same recipe used by Santos et al. (2000) for CORALIE spectra. Finally, the combined high S/N HARPS HD 219828 was observed 22 times with the HARPS spectro- spectra were analyzed to derive stellar atmospheric parameters graph, between May 2005 to August 2006. Each measurement and iron abundances using the method described in Santos et al. was done using an exposure time of 900 s, in order to average out (2004b). the stellar oscillation noise (e.g Bouchy et al. 2005). The com- plete radial velocity measurements obtained and the errors are presented in Table 2. It is worth noticing that the errors quoted 3. A new planet around HD 219828 in Table 2, which are used to plot the error bars, refer solely to 3.1. Stellar characteristics the instrumental (calibration) and photon-noise error share of the total error budget (e.g. activity and/or stellar oscillations are not According to the Hipparcos catalog (ESA 1997), HD 219828 is a considered, given the difficulty in having a clear estimate of their G0IV star with a parallax π = 12.33±1.01mas, an apparentmag- influence). nitude mv = 8.04, and a colour index B − V = 0.654. From these A quick analysis of the data revealed the presence of a stable parameters, and taking the bolometric correction from Flower 3.83-day period radial-velocity signal. This signal is best fitted −1 (1996), we derive an absolute magnitude Mv = 3.49 and a lumi- using a Keplerian fit with an amplitude K of 7.0 m s , and a nosity of 3.34 L(. non-significant eccentricity (see Fig. 1; leaving the eccentricity The analysis of our combined HARPS spectra, with a total as a free parameter in the orbital fit provides a value of 0.11 ± S/N ratio above 500, provide Teff = 5891 ± 18, log g = 4.19 ± 0.09). This corresponds to the expected signal induced by the 0.02, and [Fe/H] = 0.19±0.02. Along with the above mentioned presence of a 19.8 Earth-mass (minimum-mass) companion to absolute magnitude, the Geneva Evolutionary models (Schaerer HD 219828 (Table 3). et al. 1993) indicate that HD 219828 has a mass of 1.24 M( and Our radial velocity measurements as a function of Julian an age of 5.2 Gyr. In Table 1 we summarize the stellar parameters Date are shown in Fig. 2. The plot shows that, in addition derived for this star. to the short period signal, there is a long-period trend super- HD 219828 shows a slightly lower spectroscopic gravity as imposed. We have tried to fit a two Keplerian model to the compared to main-sequence dwarf stars of the same temperature. whole radial-velocity set using the stakanof genetic algorithm, An even lower gravity (log g = 4.03) is obtained if the isochrone recently developed by Tamuz et al. (in preparation; see also

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Table 2. Radial-velocity measurements of HD 219828 and HD 102195.

−1 −1 JD Vr [km s ] σ(Vr) [km s ] HD 219828 2 453 509.928056 –24.03256 0.00066 2 453 510.928367 –24.01790 0.00091 2 453 701.542843 –24.02993 0.00041 2 453 702.532514 –24.01817 0.00046 2 453 703.531084 –24.01824 0.00036 2 453 704.530683 –24.02639 0.00040 2 453 705.544423 –24.02791 0.00035 2 453 706.530181 –24.01890 0.00045 2 453 707.529860 –24.02140 0.00059 2 453 708.551863 –24.03220 0.00046 2 453 709.539910 –24.02877 0.00057 2 453 710.539622 –24.01823 0.00049 2 453 930.824911 –24.05650 0.00066 2 453 931.779782 –24.05201 0.00104 2 453 932.776323 –24.04571 0.00065 2 453 933.735140 –24.04473 0.00080 2 453 934.739643 –24.05830 0.00070 2 453 935.739712 –24.05720 0.00069 2 453 936.730288 –24.04521 0.00076 2 453 946.762931 –24.06075 0.00071 2 453 951.833961 –24.05168 0.00074 2 453 975.734459 –24.05942 0.00069 Fig. 1. Top: phase-folded radial-velocity measurements of HD219828, HD 102195 and the best Keplerian fit to the data with a period of 3.8-days. In this fit 2 453 501.574413 2.15457 0.00037 the long term radial-velocity trend was subtracted. See text for more de- 2 453 503.580640 2.06803 0.00133 tails. Bottom: radial-velocity measurements of HD 219828 as a function 2 453 504.587954 2.15706 0.00082 of time for one epoch of measurements. 2 453 506.621511 2.08172 0.00146 2 453 550.535485 2.17258 0.00075 2 453 551.523110 2.11355 0.00117 Table 3. Elements of the fitted orbit for HD 219828b. 2 453 757.789800 2.06038 0.00038 2 453 761.827482 2.05137 0.00036 P 3.8335 ± 0.0013 d 2 453 765.771759 2.07657 0.00037 T 2453898.6289 ± 0.072 d 2 453 785.772065 2.13058 0.00037 e 0.0† 2 453 788.796126 2.16390 0.00033 V −24.032 ± 0.001 km s−1 2 453 927.478651 2.08383 0.00062 r ω 0.0† degr 2 453 930.461429 2.05648 0.00062 K 7.0 ± 0.5 m s−1 2 453 931.470554 2.06908 0.00067 1 f (m)0.1364 × 10−12 ± 0.3572 × 10−12 M 2 453 932.485026 2.17049 0.00065 1 ( σ(O−C) 1.7 m s−1 2 453 933.465358 2.18322 0.00069 N 22 2 453 934.470134 2.08446 0.00062 m sin i 19.8 M 2 453 935.471600 2.06340 0.00065 2 ⊕ 2 453 936.476701 2.13495 0.00061 † Fixed to zero. The obtained eccentricity was consistent with a circular orbit according to the Lucy & Sweeney (1971) test.

Pepe et al. 2007). Although the 3.83-day period signal was al- ways present in the solutions, several similar quality solutions were found for the long period signal. These always ranged from ∼180 to ∼800 days in period, corresponding to the presence of a To understand if the short period radial-velocity signal ob- planet in the Jupiter-mass domain (∼0.7 MJup). The best fit (with served could be due to the presence of stellar spots (e.g. Saar an rms of only 1.2 m s−1) gave a period of 181-days, eccentricity & Donahue 1997; Queloz et al. 2001) or stellar blends (Santos of 0.3, and amplitude K = 21.6 m s−1 for the long period orbit. et al. 2002), we computed the Bisector Inverse Slope (BIS) of Given the ambiguity in the result, we prefer at this point to fit the HARPS Cross-Correlation Functions (CCF) of HD 219828. a simple quadratic drift to the residuals of the short period sig- In Fig. 3 we plot the resulting BIS as a function of the radial- nal. The global rms of the Keplerian + drift solution is 1.7 m s−1. velocity, after having subtracted the long period quadratic trend We note, however, that this long term trend does not perfectly fit to the data. The result shows that no correlation exists be- the observed variation; some of the two Keplerian models had tween the two quantities, suggesting that stellar activity of stellar better χ2 than the adopted preliminary solution. This is also il- blends cannot explain the short period and low amplitude radial- lustrated by the increasing residuals of the fit as a function of velocity variation observed. time (Fig. 2). These facts explain the relatively high rms of the Together with the very low activity level of the star, we thus orbital solution found, and may imply that the orbital parame- conclude that the 3.83-day orbital period observed can be bet- ters (e.g. e and T) of the short period planet may change slightly ter explained by the presence of a Neptune-mass planet orbiting when the long period solution is better constrained. HD 219828.

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Table 4. Stellar parameters for HD 102195.

Parameter Value Reference Spectral type K0V Hipparcos Parallax [mas] 34.51 ± 1.16 Hipparcos Distance [pc] 28.98 Hipparcos mv 8.07 Hipparcos B − V 0.835 Hipparcos Mv 5.76 – Luminosity [L(] 0.47† – Mass [M(] 0.87 Schaller et al. (1992) & log RHK −4.56 HARPS v sin i [km s−1] 2.6†† Teff [K] 5291 ± 34 log g 4.45 ± 0.04 ξt 0.89 ± 0.05 [Fe/H] 0.05 ± 0.05 † Using the bolometric correction of Flower (1996). †† From HARPS spectra using a calibration similar to the one presented by Santos et al. (2002).

Fig. 2. Top: radial-velocity measurements of HD219828 as a function of time, and a fit to the data including a 3.8-day period Keplerian and a using the temperature, astrometric distance, and stellar luminos- long period quadratic trend. Bottom: residuals of the fit. ity (log g = 4.56). A stellar mass of 0.87 M( is derived from the absolute magnitude, effective temperature, and metallicity, after comparison with stellar evolution models (Schaller et al. 1992). The stellar radius estimated using the luminosity–temperature– radius relation is 0.82 R(. According to the properties summa- rized in Table 4, HD 102195 is a typical solar metallicity main- sequence dwarf. The HARPS cross-correlation function gives an estimate of the projected stellar rotational velocity v sin i = 2.6 km s−1,a value that is reasonably compatible with the one derived by (Ge et al. 2006). Using HARPS spectra, we further derive a rather high value for the chromospheric activity index for this star & = − & = − (log RHK 4.56). A similarly high value of log RHK 4.30 was obtained by Strassmeier et al. (2000). Using the derived chromospheric activity index, and the colour of the star, we used the Henry et al. (1996) calibration to derive an estimate for the stellar age of 1.17 Gyr, and a rotational period of about 20-days (Noyes et al. 1984), slightly above the value found using the photometric light curve (∼12-days – Ge et al. 2006). As an additional age indicator, the Li abundance was computed as described in Israelian et al. (2004). We found a Li abundance log ( (Li) < 0.2. According to Sestito & Randich (2005), a Teff ∼ 5300 K star with its lithium content should have Fig. 3. BIS vs. radial-velocity for HD 219828. The best linear fit to the an age above 0.5 Gyr, also compatible with the estimate of Ge data is shown, and has a slope compatible with zero. et al. (2006).

4. The planet around HD 102195 4.2. The HARPS orbital solution 4.1. Stellar characteristics of HD 102195 The discovery of a giant planet around HD 102195 has been re- cently announced by the Exoplanet Tracker team1 using a new According to the Hipparcos catalog (ESA 1997), HD 102195 is a concept of instrument combining spectroscopy and interferome- K0V star with visual magnitude mv = 8.07 and colour index B − try (Ge et al. 2006). Since this star is among our surveyed sam- V = 0.835. From the visual magnitude and astrometric distance ple, it is interesting to compare the results obtained by Ge et al. of 28.98 pc (π = 34.51 ± 1.16), we derive an absolute magnitude (2006) with those obtained with HARPS. Mv = 5.76. Considering the bolometric correction taken from We observed HD 102195 four times in May 2005, in the Flower (1996), the obtained luminosity is then 0.47 L(. framework of our program aimed at finding earth-mass plan- Accurate stellar atmospheric parameters for HD 102195 ets with HARPS around a metallicity biased sample of stars. were derived from the analysis of our combined HARPS spec- Based on these first four measurements, HD 102195 was clearly tra, with a S/N ratio above 300. Our analysis provides Teff = detected to be a radial velocity variable. These first four mea- 5291 ± 34 K, log g = 4.45 ± 0.04, and [Fe/H] = 0.05 ± 0.05, surements did not show any correlation between the bisector and in excellent agreement with previous estimates (Ge et al. 2006). This value for the gravity is compatible with the one derived 1 http://www.astro.ufl.edu/et/

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Table 5. Elements of the fitted orbit for HD 102195b, using the original and the corrected radial-velocity measurements. See text for more details.

Original velocities Corrected velocities P 4.114225 ± 0.000855 4.113775 ± 0.000557 d T 2453896.00 ± 0.04 2453895.96 ± 0.03 d e 0.0† 0.0† −1 Vr 2.113 ± 0.002 2.132 ± 0.001 km s ω 0.0† 0.0† degr −1 K1 64 ± 3 63 ± 2 m s −10 −10 −10 −10 f1(m)1.135 × 10 ± 0.169 × 10 1.070 × 10 ± 0.103 × 10 M( σ(O−C) 9.4 6.1 m s−1 N 19 19 m2 sin i 0.46 0.45 MJup † Fixed to zero. The obtained eccentricity was consistent with a circular orbit according to the Lucy & Sweeney (1971) test.

Fig. 5. BIS vs. radial-velocity for HD 102195. The best linear fit to the data is shown, and has a slope compatible with zero. Fig. 4. Top: phase-folded radial-velocity measurements of HD 102195, and the best Keplerian fit to the data with a period of 4.1-days. Bottom: residuals of the Keplerian fit as a function of time. by stellar spots or stellar blends, and supports the planetary explanation. Using the best-fit orbital solution and the mass of the radial velocities, suggesting that this star could be indeed a HD 102195, we derive for the companion a minimum planet host. mass 0.46 MJup, in close agreement with the 0.49 MJup found The HARPS GTO team was contacted and kindly accepted by Ge et al. (2006). to monitor HD 102195. In total, and counting also the data points The stellar radius together with the rotational period men- obtained in our subsequent observing runs, 19 radial-velocity tioned above provide a value for the stellar rotational veloc- points were added between June 2005 and February 2006 in or- ity of ∼3.5 km s−1. Comparing this with the estimated value for der to derive a good orbital solution. In the meanwhile, the star the projected rotational velocity (v sin i = 2.6 km s−1) we get was announcedas a planethost by Geet al.(2006).In Table 2 we sin i = 0.74, corresponding to an inclination of 47 degrees. If the show the radial velocity measurements collected for HD 102195 orbital axis of the planetary orbit is aligned with the stellar ro- between May 2005 and February 2006 covering a time span of tation axis (e.g. Queloz et al. 2000), such an inclination could about 300 days. explain why no transit of HD 102195 b was detected by Ge et al. Using the radial velocity data shown in Table 2 an orbital (2006). The real mass of planet would then be 0.62 MJup. solution was found (Table 5). The best fit Keplerian orbit gives a period of 4.11-days, an amplitude K of 64 m s−1, and a non- 4.3. Correcting the radial velocities for the stellar noise significant eccentricity. These values are similar to the ones de- rived by Ge et al. (2006). The phase-folded radial velocity plot In the case of HD 102195, the residuals (O−C) given in Table 5 of the measurements is shown in Fig. 4. of 9.4 m s−1 are much higher than the errors on radial velocity Despite the high activity level of the star, an analysis of quoted in Table 2. This extra-noise is likely due to stellar activ- the bisector of the HARPS cross-correlation functions shows no ity (see previous section). Ge et al. (2006) found that this star clear correlation with the radial-velocities (Fig. 5). This strongly has a photometric variation with a period of ∼12 days, and an suggests that the radial-velocity signal is not being produced amplitude of ∼0.015 mag.

97 Chapter 4. Main results

726 C. Melo et al.: A new Neptune-mass planet orbiting HD 219828

Fig. 6. BIS as a function of the residuals of the 4.1-day Keplerian radial- Fig. 7. Transit depth expected for three different planet models as a velocity fit for HD 102195. The data is clearly correlated. The best lin- function of the mass. The lower curve represents the model of Valencia ear fit has a slope of −0.73. et al. (2006), and the two upper curves the scaled radii considering the mean density of Neptune and the icy Saturn satellite Rhea. The vertical dotted line represents the location of HD 219828 b. The signature of its high activity level is also seen on the bisector analysis done in this paper. Figure 6 shows that the residuals of the best Keplerian orbit (previous section) do cor- of a m2 sin i = 0.45 MJup planet orbiting HD 102195, previously relate with the bisector measurements in a clear indication that announced by another team. the high-activity level of this star is responsible for the observed The planet around HD 219828 is the 11th found with masses high residuals (e.g. Saar & Donahue 1997; Santos et al. 2000). similar to the mass of Neptune. As the number of these sys- According to Fig. 6, the relation between the residual radial- tems increases, new statistical analysis will be possible. In this velocities and the BIS measurements can be approximated by a sense, Udry et al. (2006) have discussed the possibility that the linear function a×BIS +b whose coefficients were determined by well known strong correlation between the presence of planets a least-squares fit. This function can be subtracted from the ra- and the stellar metallicity that exists for stars hosting giant plan- dial velocity data points in order to correct for the stellar activity. ets (e.g. Gonzalez 1997; Santos et al. 2001, 2004b; Fischer & Using those corrected velocities, a new orbital solution was Valenti 2005) does not seem to be present for their lower mass found. The results, also presented in Table 5, show a clear im- counterparts. For very-low mass companions, the metallicity dis- provement. Not only the rms around the orbital fit was reduced tribution could rather be flat, something that may be explained by from 9.4 to 6.1 m s−1 (equivalent to subtracting quadratically a recent models of planet formation (Ida & Lin 2004; Benz et al. noise of 7 m s−1), but also the estimated errors in all orbital pa- 2006). Although the metallicity of HD 219828 is rather high, we rameters decreased by almost a factor of two. The resulting mass must be careful when analyzing the numbers, since the planet- search project described in this paper concentrated its efforts in for HD 102195 b also slightly decreased (0.45 MJup) due to the small decrease in the amplitude of the orbital solution. studying a sample of metal-rich stars only. This result shows that the study of the BIS can be used to An important question triggered by the previous discover- correct, at least to some extent, the noise on the radial-velocities ies is whether these small planets are constituted by rocks, ice introduced by intrinsic stellar activity. If the BIS vs. RV corre- or a mixture of both. A more solid assessment of the nature of lation can be calibrated (observationally or theoretically) as a HD 219828 b can only be given if the information about its mass function of the stellar properties, we could think that in the fu- and radius is available. The detection of a transiting signature ture an externally determined BIS vs. RV function can be used would give us information about the radius and mean density of to correct the radial velocity data. This will be extremely impor- the planet, thus opening the possibility to unveil its composition, tant in reaching accuracy of the other of the cm s−1 as expected like is the case for the giant planets found to transit their host in future instruments like CODEX (Pasquini et al. 2005). stars (e.g. Konacki et al. 2003; Pont et al. 2004). Although the short period of the orbit implies a reasonable value for the transit probability (∼15%), the measure of a tran- 5. Concluding remarks sit of HD 219828 b is not an easy task. Using the models pre- sented by Valencia et al. (2006), considering that this m2 sin i = In this paper, we present the discovery of a new 20 M⊕ 19.8 M⊕ planet could have an Earth-like composition, we derive (minimum-mass) short period planet, orbiting the sun-like star its radius to be ∼2.2 Earth radii. Although this value must be HD219828. The presence of a second longer period compan- taken as a lower limit (both due to the lower limit of the mass, ion to the system, likely a Jovian planet, is also clear from and to the fact that it may have an extended atmosphere – Lovis our data. Unfortunately, the bad phase coverage of the radial- et al. 2006), observing the transit of such a low-mass planet velocity points does not permit to settle the orbital parame- is likely impossible with current ground-based instrumentation ters of this latter companion. Finally, we confirm the existence (see Fig. 7).

98 4.4. Additional discoveries

C. Melo et al.: A new Neptune-mass planet orbiting HD 219828 727

For the case of HD 102195, we have also tested, for the first Gonzalez, G. 1997, MNRAS, 285, 403 time, the use of the bisector for the cross-correlation function Henry, T. J., Soderblom, D. R., Donahue, R. A., & Baliunas, S. L. 1996, AJ, 111, to correct the radial-velocity measurements for the noise intro- 439 Ida, S., & Lin, D. N. C. 2004, ApJ, 616, 567 duced by stellar activity. The results show that a considerable Israelian, G., Santos, N. C., Mayor, M., & Rebolo, R. 2004, A&A, 414, 601 improvement can be found. Although the technique may proba- Konacki, M., Torres, G., Jha, S., & Sasselov, D. 2003, Nature, 421, 507 bly be improved, this conclusion constitutes a first step towards Lovis, C., Mayor, M., Pepe, F., et al. 2006, Nature, 441, 305 the subtraction of the radial-velocity noise produced by stellar Lucy, L. B., & Sweeney, M. A. 1971, AJ, 76, 544 McArthur, B. E., Endl, M., Cochran, W. D., et al. 2004, ApJ, 614, L81 activity phenomena. Noyes, R. W., Hartmann, L. W., Baliunas, S. L., Duncan, D. K., & Vaughan, A. H. 1984, ApJ, 279, 763 Acknowledgements. We thank the referee Dr. Barbara McArthur whose com- Pace, G., & Pasquini, L. 2004, A&A, 426, 1021 ments helped to improve our paper. Support from Fundação para a Ciência e Pasquini, L., Cristiani, S., Dekker, H., et al. 2005, The Messenger, 122, 10 a Tecnologia (Portugal) to N.C.S. and S.G.S. in the form of fellowships (refer- Pepe, F., Mayor, M., Queloz, D., et al. 2004, A&A, 423, 385 ences SFRH/BPD/8116/2002 and SFRH/BD/17952/2004) and a grant (reference Pepe, F., Correia, A. C. M., Mayor, M., et al. 2007, A&A, 462, 769 POCI/CTE-AST/56453/2004) is gratefully acknowledged. WG, GP and MTR Pont, F., Bouchy, F., Queloz, D., et al. 2004, A&A, 426, L15 were supported by the Chilean FONDAP Center of Astrophysics 15010003. Queloz, D., Eggenberger, A., Mayor, M., et al. 2000, A&A, 359, L13 Queloz, D., Henry, G. W., Sivan, J. P., et al. 2001, A&A, 379, 279 Rivera, E. J., Lissauer, J. J., Butler, R. P., et al. 2005, ApJ, 634, 625 References Saar, S. H., & Donahue, R. A. 1997, ApJ, 485, 319 Santos, N. C., Mayor, M., Naef, D., et al. 2000, A&A, 361, 265 Alibert, Y., Baraffe, I., Benz, W., et al. 2006, A&A, 455, L25 Santos, N. C., Israelian, G., & Mayor, M. 2001, A&A, 373, 1019 Benz, W., Mordasini, C., Alibert, Y., & Naef, D. 2006, in Tenth Anniversary Santos, N. C., Mayor, M., Naef, D., et al. 2002, A&A, 392, 215 of 51 Peg-b: Status of and prospects for hot Jupiter studies, ed. L. Arnold, Santos, N. C., Bouchy, F., Mayor, M., et al. 2004a, A&A, 426, L19 F. Bouchy, & C. Moutou, 24 Santos, N. C., Israelian, G., & Mayor, M. 2004b, A&A, 415, 1153 Bonfils, X., Udry, S., Delfosse, X., Santos, N., & Mayor, M. 2005, A&A, 443, Schaerer, D., Charbonnel, C., Meynet, G., Maeder, A., & Schaller, G. 1993, L15 A&AS, 102, 339 Bouchy, F., Bazot, M., Santos, N. C., Vauclair, S., & Sosnowska, D. 2005, A&A, Schaller, G., Schaerer, D., Meynet, G., & Maeder, A. 1992, A&AS, 96, 269 440, 609 Sestito, P., & Randich, S. 2005, A&A, 442, 615 Butler, R. P., Vogt, S. S., Marcy, G. W., et al. 2004, ApJ, 617, 580 Strassmeier, K., Washuettl, A., Granzer, T., Scheck, M., & Weber, M. 2000, ESA. 1997, The Hipparcos and Tycho Catalogues A&AS, 142, 275 Fischer, D. A., & Valenti, J. 2005, ApJ, 622, 1102 Udry, S., Mayor, M., Benz, W., et al. 2006, A&A, 447, 361 Flower, P. J. 1996, ApJ, 469, 355 Valencia, D., O’Connell, R. J., & Sasselov, D. 2006, Icarus, 181, 545 Ge, J., van Eyken, J., Mahadevan, S., et al. 2006, ApJ, 648, 683 Vogt, S. S., Butler, R. P., Marcy, G. W., et al. 2005, ApJ, 632, 638

99 Chapter 4. Main results

100 4.5. Radial-velocity variable stars with no detected planet

4.5 Radial-velocity variable stars with no detected planet

The variability of the radial velocities can be quantified using the parameter E/I, which compares external and internal dispersions of the measurements. The external dis- persion E is the standard deviation of the radial-velocity measurements (σRV ) calculated relative to the mean radial velocity. The internal dispersion I represents the average error of the individual measurements. Here, I just take the photon-noise errors into account, but I remind that other sources of errors are present (instrumental error, stellar activity). Many of the programme stars that presented a significant variability in their radial ve- locities had no orbital solution determined. As discussed in Section 2.4, there exist other sources of radial-velocity variability rather than that caused by the presence of a planetary companion. Any information concerning the age of the star, the presence of faint stellar companions, the level of chromospheric activity, or photometric variability, among others, can shed light on our understanding of the possible causes for the observed radial-velocity variation in the cases for which no reliable orbital solution could be fitted.

Below I describe some cases for which a significant variability in the radial velocities was observed but no planet could be detected. I assumed that stars having the parameter E/I greater than about 2.0 are possible radial-velocity variables. Some of them are classified in the literature as being photometric variable, having a strong level of chromospheric activity, and/or being member of a multiple system. Others present no evidence of photometric variability, chromospheric activity, or duplicity. Some of them have only the information of being member of an young moving group, which could indicate a high level of activity.

For the stars listed in Subsections 4.5.1 and 4.5.2 for which no value of the indices & log RHK or S is available in the literature to quantify their chromospheric activity level, I present in Figures 4.7 and 4.9 their summed spectra obtained with ELODIE in the region of the λ3968.5 Ca II H absorption line, in order to illustrate an eventual emission (or the absence of emission) in the core of such a line.

4.5.1 Photometric variable and/or active stars

Panels (a) to (h) of Figure 4.6 present the radial-velocity measurements of stars for which the available data in the literature indicates some photometric variability and/or a high level of chromospheric activity. Some of the physical parameters of each star are listed as well. They are: the metallicity [Fe/H], the projected rotation velocity v sin i, the amplitude of the radial-velocity variation ∆RV, its standard deviation σRV , the total number of days from first to last measurement ∆JD, and the external-to-internal dispersion ratio E/I.

HD 691 (Figure 4.6a) This star is present in the list of variable stars of Kazarovets et al. (1999), and was classified as having a possible variability type related to the stellar

101 Chapter 4. Main results

chromospheric activity and rotating spots. The amplitude of variation is of 0.06 mag in the Hp-band, which is the magnitude in the system of the Hipparcos catalogue. Indeed, some indicators of the chromospheric activity level label it as a very active & star: Strassmeier et al. (2000) measured a index logRHK= 4.15, whereas Wright & − et al. (2004) derived log R = 4.38 and S =0.446. Based on its space velocities HK − (U, V, W ), Montes et al. (2001) considered that this is a young disk star, which supports the derived high activity level.

HD 24053 (Figure 4.6b) Strassmeier et al. (2000) determined a significantly high activity level for this star, measuring a index log R& = 4.30. HK − HD 28099 (Figure 4.6c) This is a very young star member of the Hyades moving group, which is aged 625 Myr. The index of chromospheric activity obtained by Paulson et & al. (2002) is log RHK= 4.47. Lyra & Porto de Mello (2005), using the Hα absorption − & line instead of the Ca II H & K lines, determined log F = 4.14. In addition to Hα − the high level of chromospheric activity, Guenther et al. (2005) detected a M dwarf companion in a separation of 0.43&&. The apparent magnitude of the primary in the & K -band is mK# = 6.6, whereas the secondary has mK# = 11.97. HD 30572 (Figure 4.6d) The Hipparcos epoch photometry reveals a variation of more than 0.1 mag in the Hp-band. The first radial-velocity measurements we made show a significant variation of more than 100 m s−1 in a few days, but the last ones seems to be rather constant. The observed variation, both in photometry and in radial velocity, could be the result of chromospheric activity. No index of activity is available in the literature, but our spectra gathered with ELODIE indicate some emission in the core of the Ca II H line (see Figure 4.7).

HD 31000 (Figure 4.6e) Member of the Ursa Major moving group (Montes et al. 2001), with an age of 300 Myr, this is also an active star, as indicated by the indices log R& = 4.10 (Strassmeier et al. 2000) and logR& = 4.419 (King et al. 2003). HK − HK − Moreover, it is presented in the catalogue of variable stars of Kazarovets et al. (2006), who derived an amplitude of variability of 0.03 mag in the V -band induced by the chromospheric activity.

HD 82443 (Figure 4.6f) According to Montes et al. (2001), this star is a young member of the local association, which is aged 20-150 Myr. Its level of chromospheric activity is very high: Strassmeier et al. (2000) determined an index logR& = 4.05, Duncan HK − et al. (1991) measured S =0.674, and Young et al. (1989) obtained S =0.648. The Hipparcos epoch photometry indicates an amplitude of variation of more than 0.01 mag in the Hp-band. Furthermore, it is classified in the Simbad database as having a periodic photometric variability. Indeed, many authors found a photometric variation with a period of about P 5 days that they attributed to cold photospheric ∼ spots across the stellar disk. Henry et al. (1995) determined a rotation period of 5.43 days, with an amplitude that changes from year to year between 0.03 and 0.08 mag. Messina et al. (1999) and Messina & Guinan (2003) were able to estimate that the rotation period is not constant, changing from 5.25 to 5.48 days ∼ ∼

102 4.5. Radial-velocity variable stars with no detected planet

(a) (b)

(c) (d)

(e) (f)

Figure 4.6: Stars having a significant radial-velocity variation (E/I ! 2.0) and for which some available data in the literature classify them as photometric variables and/or chro- mospheric actives.

103 Chapter 4. Main results

(g) (h)

Figure 4.6: Continued from previous page.

during the stellar cycle of activity. This star was also observed with the SOPHIE spectrograph. As in the case of the ELODIE measurements, the radial velocities show a large amplitude of variation (about 300 m s−1). However, no orbital solution could be fitted to the data. The line bisectors show a variation of about 200 m s−1, but no correlation with the radial velocities nor with the stellar rotation period of 5 days is observed.

HD 108575 B (Figure 4.6g) Classified in the Simbad database as a star in double system. Montes et al. (2001) listed this star as a member of the IC 2391 supercluster, aged 35-55 Myr. Zuckerman et al. (2006) suggested that it is a possible member of the Carina-near moving group, which is aged about 200 Myr. Its age indicates a strong & level of chromospheric activity, which is indeed confirmed by the index log RHK= 4.22 (Strassmeier et al. 2000). Moreover, it presents a strong photometric variation − of more than 0.7 mag in the Hp-band of the Hipparcos epoch photometry.

HD 121979 (Figure 4.6h) Kazarovets et al. (2006) classified this star as variable. The amplitude of variability is about 0.03 mag in the V -band, and is related to rotating spots and chromospheric activity. Koen & Eyer (2002) determined P 5 days for ∼ the period of the photometric variation. It is also a member of the young IC 2391 supercluster, as suggested by Montes et al. (2001). The level of activity is very strong as indicates the index log R& = 4.22 from Strassmeier et al. (2000). HK −

104 4.5. Radial-velocity variable stars with no detected planet

Figure 4.7: Summed ELODIE spectra in the region of the Ca II absorption line at λ3968.5 for the star HD 30572. Despite the low spectral quality, we can observe some emission in the core of the line. For clarity, spectral features due to pollution of thorium emission lines were removed.

105 Chapter 4. Main results

4.5.2 Some peculiar cases

Panels (a) to (j) of Figure 4.8 present the radial-velocity measurements of stars for which the available data in the literature present no clear information that could explain the observed variation in the radial velocities. As in Figure 4.6, some of the physical parameters of each star are also listed in the panels.

HD 3141 (Figure 4.8a) This star is not chromosphericaly active as indicates the index log R& = 5.00 derived by Strassmeier et al. (2000), and logR& = 4.97 and S = HK − HK − 0.187 determined by Wright et al. (2004).

HD 7352 (Figure 4.8b) The only relevant information available for this star is that there is no evidence of duplicity, which is the conclusion of Horch et al. (2002) based on high-quality observations. Figure 4.9a shows the averaged spectrum of our ELODIE measurements in the region of the Ca II H line, and no emission is observed in the core of the line.

HD 29836 (Figure 4.8c) According to the indices logR& = 5.13 and S =0.142 deter- HK − mined by Wright et al. (2004), this star is not active. In addition, no photometric variability is detected in the Hipparcos data. Bonavita & Desidera (2007) classified this star as a member of a triple system, but the companions have wide separations.

HD 30589 (Figure 4.8d) This star is also a very young Hyades member. Some indicators of chromospheric activity show that it is situated in the limit between active and non-active stars: log R& = 4.82 (Paulson et al. 2002) and S =0.201 (Duncan et HK − al. 1991). Nevertheless, the last radial-velocity measurements indicate no significant variation. Some measurements made using SOPHIE show no important variation.

HD 75528 (Figure 4.8e) This star is not chromosphericaly active, as indicated by the indices S =0.152 from Young et al. (1989) and S =0.149 from Middelkoop (1982). However, it is classified by Hartkopf & McAlister (1984) as an occultation binary, with a vector separation of 0.05&&. According to Nordstr¨om et al. (2004), it is aged 5.2 Gyr, and the magnitude difference to ZAMS is 1.96 mag. Three measurements were made using SOPHIE, and no short-term variation is present.

HD 80869 (Figure 4.8f) Aged 5.4 Gyr and with a mass of 1.19 M! (Nordstr¨om et al. 2004), such a star has a magnitude difference to ZAMS of 1.36 mag, indicating a slight degree of evolution. The averaged spectrum of our ELODIE measurements in the region of the Ca II H line is somewhat noisy but, nevertheless, no emission seems to be present in the core of the line (see Figure 4.9b).

HD 86460 (Figure 4.8g) In spite of the significant radial-velocity variation (> 160 m s−1 pic-to-pic), no periodic signal could be found in our measurements. The averaged spectrum in the region of the Ca II H line is very noisy, and does not allow us to observe if any chromospheric emission is present or not.

106 4.5. Radial-velocity variable stars with no detected planet

HD 100796 (Figure 4.8h) The Hipparcos epoch photometry indicates an amplitude of variation of more than 0.1 mag in the Hp-band. In addition, Nordstr¨om et al. (2004) determined a magnitude difference to ZAMS of 1.42 mag. This star seems to be not chromospherically active as indicated by the absence of emission in the core of the Ca II H line (see Figure 4.9c).

HD 147187 (Figure 4.8i) Besides the Hipparcos epoch photometry, which shows no sig- nificant variability, this star has no relevant information available in the literature. The region of the Ca II H line is plotted in the Figure 4.9d, and shows no emission in the line core.

HD 163589 (Figure 4.8j) A trend seems to be present in the radial velocities of this star, with an amplitude of variation of about 50 m s−1 in one year. It is just a bit older than the Sun (6.7 Gyr), and have a mass slightly larger (1.14 M!). On the other hand, the magnitude difference to ZAMS derived by Nordstr¨om et al. (2004) is 1.88 mag, indicating a rather significant degree of evolution. Figure 4.9e shows the averaged spectrum of our ELODIE measurements in the region of the CaII H line, and no emission in the core of the line is observed.

107 Chapter 4. Main results

(a) (b)

(c) (d)

(e) (f)

Figure 4.8: Stars having a significant radial-velocity variation (E/I ! 2.0) but for which the available data in literature show no clear information that could explain the observed variability.

108 4.5. Radial-velocity variable stars with no detected planet

(g) (h)

(i) (j)

Figure 4.8: Continued from previous page.

109 Chapter 4. Main results

(a) (b)

(c) (d)

(e)

Figure 4.9: Summed ELODIE spectra in the region of the Ca II absorption line at λ3968.5. They all show no emission in the core of the line. For clarity, spectral features due to pollution of thorium emission lines were removed.

110 4.5. Radial-velocity variable stars with no detected planet

4.5.3 Binary or brown dwarf stars

Some of the observed target stars present a huge radial-velocity variation, and re- vealed to be either members of a binary system, or to have an unseen faint companion, like a brown dwarf star. They were monitored either with the ELODIE or the SOPHIE spectrograph, and some of them have also a few measurements made with CORAVEL (the predecessor of ELODIE). Those monitored with SOPHIE have also one measurement made using ELODIE (except the star HD 10790), which allowed an estimate of metallicity and projected rotation velocity with the calibrations described in Section 3.3. Similar calibrations are not yet available for stars observed using SOPHIE. Panels (a) to (n) of Figure 4.10 present the radial-velocity measurements of such stars together with some relevant physical parameters. Panels (a) to (d) of Figure 4.11 display the cases of stars for which a Keplerian orbital solution could be fitted to the available data from CORAVEL and ELODIE. The orbital parameters are also listed, including the difference between the velocity zero-point of the two datasets (∆C−E). Below I give a brief description of each one of these stars.

HD 10790 (Figure 4.10a) The number of measurements available for this star is not enough to allow the fit of any reliable orbital solution. Given the radial-velocity variation of CORAVEL together with SOPHIE observations, I could only estimate a lower limit for the mass of the companion, namely M2 sin i>30 MJup.

HD 14651 (Figure 4.10b) Combining both CORAVEL and SOPHIE radial-velocities, a period of variation of about 100 days seems to be present. However, the number of measurements is still small, and no reliable orbital solution could be found. I estimated a lower limit of 40 MJup for the mass of the companion.

HD 15292 (Figure 4.10c) Taking the CORAVEL measurements into account, I could estimate a lower limit of 150 MJup for M2 sin i, but no orbital solution. The star is probably a candidate of a binary system member.

HD 29037 (Figure 4.10d) The huge radial-velocity variation observed classify this star as a member of a binary system, but more measurements are required to provide an orbital solution. No observation was made with CORAVEL.

HD 34031 (Figure 4.10e) Again, the variation in the radial velocities is significantly large, and the star has probably a binary companion.

HD 49178 (Figure 4.10f) Probably a member of a binary system, the number of radial- velocity measurements of this star does not allow a reliable solution for the orbit.

BD+481523 (Figure 4.10g) With no measurements made with CORAVEL, the radial- velocity variation of this star could be the result of a massive planet in a short-period orbit. More measurements are required to allow better conclusions.

111 Chapter 4. Main results

(a) (b)

Figure 4.10: Binary or brown dwarf candidates for which no Keplerian orbital solution could be found. The radial-velocity measurements are from CORAVEL, ELODIE and/or SOPHIE spectrographs.

HD 78536 (Figure 4.10h) No measurement is available for this star in the CORAVEL database. The radial-velocity variation observed in the SOPHIE measurements is significantly large, and indicates that the mass of the companion is at least that of a brown dwarf star. HD 104289 (Figure 4.10i) The number of radial-velocity measurements is still small to allow an unambiguous orbital solution. Nevertheless, I could identify a period of variation of about 988 days, and an inferior limit of 65 MJup in the mass of the companion. HD 106888 (Figure 4.10j) Again, the small number of measurements of this star does not allow an unique Keplerian solution for the orbit, but a period of variation of about 595 days seems to be present, and I estimated the minimum mass M2 sin i for the companion to be between 50 and 100 MJup. HD 125193 (Figure 4.10k) A very long orbital period seems to be present. The minimum mass of the companion is probably of the order of 0.2 M! or higher. HD 161479 (Figure 4.10l) The period of variation in the radial velocities of this star seems to be short, but the number of measurements is still small to allow any orbital fit. Given the very large variation pic-to-pic, the mass of the companion is probably higher than 100 MJup. HD 199100 (Figure 4.10m) I could not fit a reliable orbital solution to the CORAVEL and ELODIE data of this star, but an orbital period of about 550 days and a minimum mass M sin i 0.4M! represent a possible and likely solution. 2 ∼ HD 212585 (Figure 4.10n) The radial velocities seems to have a very long period of variation, but more measurements are needed.

112 4.5. Radial-velocity variable stars with no detected planet

(c) (d)

(e) (f)

(g) (h)

Figure 4.10: Continued from previous page.

113 Chapter 4. Main results

(i) (j)

(k) (l)

(m) (n)

Figure 4.10: Continued from previous page.

114 4.5. Radial-velocity variable stars with no detected planet

(a)

(b)

Figure 4.11: Binary or brown dwarf candidates for which a Keplerian orbital solution was fitted to the available data. Left panels present the CORAVEL and ELODIE radial- velocity measurements, whereas right panels show our best Keplerian solution. The dis- persion around the solution is σC for CORAVEL and σE for ELODIE.

HD 62923 (Figure 4.11a) The Keplerian orbital solution fitted to the radial-velocity data of this star gives a period P = 124 days and a minimum mass M2 sin i = 0.051 M!, a probable brown dwarf candidate.

HD 89010 (Figure 4.11b) I fitted a Keplerian orbital solution to the CORAVEL and ELODIE data, which gives a period P = 788 days and a minimum mass M2 sin i = 0.274 M!. Therefore, the companion is likely a faint low mass star.

115 Chapter 4. Main results

(c)

(d)

Figure 4.11: Continued from previous page.

HD 212733 (Figure 4.11c) I fitted a Keplerian orbital solution to the CORAVEL and ELODIE data of this star, and I found that the variation in the radial velocities are due to the presence of a companion with minimum mass M2 sin i = 0.109 M! orbiting the main star with a period P = 89.9 days.

HD 216320 (Figure 4.11d) The Keplerian solution fitted to the CORAVEL and ELODIE radial-velocity measurements of this star gives a relatively short period P = 20.2 days and a minimum mass M2 sin i = 0.109 M! for the companion.

116 Chapter 5

Statistical study of the programme

Contents

5.1 Metallicity distribution ...... 117 5.2 Period distribution ...... 119 5.3 Global detectability of the programme ...... 119

In this chapter, I present a statistical study based on some physical properties of the stars observed in the context of this programme, in order to estimate the expected number of stars with planets that are included in our sample. To proceed this study, I used three different approaches: 1) in Section 5.1, I compare the metallicity distribution of the stars observed in this programme using ELODIE to the fraction of stars with planets per bin of metallicity determined by (Santos et al. 2004) on a sample of stars observed using CORALIE; 2) in Section 5.2, I compare the period distribution of the planets orbiting stars included in the CORALIE survey to the same distribution of the planets around stars present in our sample; 3) and, in Section 5.3, I show the results of Monte Carlo simulations, based on the observational data itself, made to calculate the detection level of our programme for a given planetary mass and orbital period.

5.1 Metallicity distribution

Using the ELODIE spectrograph, we observed at least once about 80% of the 1061 sample stars, and for each one we estimated a value for metallicity using the relation given by Equation 3.4. The left panel of Figure 5.1 presents the metallicity distribution of the 854 observed stars. We can verify that about 26% of the observed stars have [Fe/H] 0.1 (about 15% for 0.1 [Fe/H] < 0.2, 8% for 0.2 [Fe/H] < 0.3, and 3% for ≥ ≤ ≤

117 Chapter 5. Statistical study of the programme

Figure 5.1: Left: Metallicity distribution of the 854 observed programme stars (80% of the original sample of 1061 stars). The right-hatched area indicates the high metallicity stars, which were our main targets. Right: The dashed histogram represents the fraction of stars with planets per bin of metallicity determined by Santos et al. (2004) on a sample of 875 stars observed with CORALIE, and having at least 5 measurements. In the solid histogram, I selected only the stars harbouring planets with P<100 days and having K ! 30 m s−1.

0.3 [Fe/H] < 0.4). According to the percentage of stars with planets per metallicity bin ≤ determined by Santos et al. (2004), 10 30% of the stars with [Fe/H] 0.1 are likely to host − ≥ a giant planet: about 9, 24, and 28% respectively for the same three ranges of metallicity mentioned above (see the right panel of Figure 5.1). Applying these percentages to the 854 observed stars, we find that the number of giant planets expected to be discovered in each metallicity range is respectively 12, 17, and 7, a total of 36 planets, from which roughly 25% (9 planets) are expected to be hot Jupiters1.

The right panel of Figure 5.1 shows the distribution constructed by Santos et al. (2004) using a sample of 48 stars with planets compared to 875 stars present in the CORALIE planet-search programme and having at least 5 measurements at that time. The pro- gramme that we have initiated using the ELODIE spectrograph has a baseline of observa- tions (2 years and 5 months) that is not long enough to optimize the detection of planets in long-period orbits, especially considering that the number of measurements is still small. For this reason, it is interesting to construct the same distribution only taking into account the planets for which the orbital period is short enough to be likely detected with ELODIE in the context of our programme. Thus, I selected among the 48 planet-host stars used by Santos et al. (2004) only those having planets with P<100 days. Moreover, I only

1 About 25% of the up to date list of discovered planets have periods shorter than 10 days.

118 5.2. Period distribution

chose the stars with planets that are able to induce a semi-amplitude K ! 30 m s−1 in the radial velocities, given the fact that the precision achieved with ELODIE ( 10 to − ∼ 20 m s 1) makes difficult to detect variations that are smaller and still significant. A total of 13 stars with planets fulfil these criteria, and the distribution constructed is also shown in the right panel of Figure 5.1. We can notice that the percentage of stars with planets per metallicity bin is about 1% for 0.1 [Fe/H] < 0.2, 7% for 0.2 [Fe/H] < 0.3, and ≤ ≤ 8% for 0.3 [Fe/H] < 0.4 in the new distribution. Recalculating the number of stars with ≤ planets expected to be discovered in each one of these metallicity bins, we have about 1, 5, and 2, respectively.

In summary, I compared the distribution of the percentage of stars observed with ELODIE in the present programme, as a function of the metallicity, to the percentage of stars with planets that are likely to be found inside each bin of metallicity, based on the CORALIE observational data studied by Santos et al. (2004). I found that a total of 8 3 ± metal-rich stars are expected to be discovered in our programme harbouring a giant planet with period smaller than 100 days and able to induce a semi-amplitude of variation larger than 30 m s−1 in the radial velocities.

5.2 Period distribution

Another approach to verify what is the number of stars harbouring a giant planet that we are able to detect with our programme is to compare the number of planets as a function of the orbital period orbiting stars included in the CORALIE programme to the number of planets effectively found orbiting stars present in our sample. Figure 5.2 shows the distribution of the number of planets per period range. Again, I only consider stars with [Fe/H] 0.1, and with planets inducing a semi-amplitude of variation larger − ≥ than 30 m s 1 in the radial velocities. We can notice in this figure that 7 stars present in the CORALIE sample have planets with P<100 days and satisfy the latter conditions, whereas in our sample we found 6 of such host stars.

For periods longer than 100 days, the distributions of Figure 5.2 shows that the number of stars with planets in the CORALIE programme are significantly larger than those we found in this programme, which is just the consequence of the long baseline of the CORALIE observations (almost 10 years) and the large number of measurements (at least 10 for most of the observed stars).

5.3 Global detectability of the programme

Another way to evaluate the capabilities of our metallicity biased search for hot Jupiters using the ELODIE spectrograph is to perform Monte Carlo simulations based on the

119 Chapter 5. Statistical study of the programme

Figure 5.2: Distribution of the number of planets as a function of the orbital period that were discovered in stars included in the CORALIE planet-search programme (solid histogram) compared to the same distribution but for stars present in our sample (hatched histogram). Only stars with [Fe/H] 0.1 and harbouring planets that are able to induce − ≥ a semi-amplitude K>30 m s 1 are represented. observational data itself collected in the course of the programme. Therefore, I performed numerical simulations to estimate the detection level of our programme in different points in a (P,Mp) grid, as well as to estimate the detection limits in the planetary mass for a given orbital period, taking into account: 1) the stellar mass; 2) the date of the radial- velocity measurements; 3) and the errors in each individual radial-velocity value. The date of the measurements corresponds to the Julian date of all individual radial-velocity measurements of a given star made in the course of our observation programme.

I estimated the stellar masses by doing a linear interpolation in the Table 19.18 of Cox (2000), which lists a mass-luminosity relation (absolute magnitude at the V band as a function of mass) derived for solar main-sequence stars. I used the absolute magnitude values listed in Table A.1, which are derived from the Hipparcos catalogue data. Actually, this mass-luminosity relation does not take the stellar metallicity into account. In order to verify what is the impact of the metallicity on the mass determination, I compared the mass of some stars derived using such a relation to those derived using the Geneva models of stellar evolution (Schaerer et al. 1993), with the effective temperature derived from Flower (1996), and the metallicity from the calibrations of Section 3.3. I found that the masses of metal-rich stars ([Fe/H] 0.2) are systematically smaller by 0.1 M! than ∼ those determined using the mass-luminosity relation of Cox (2000), which is valid for solar metallicities. I then compared the result of the simulations as follows: 1) using masses from the mass-luminosity relation; 2) using masses from the mass-luminosity relation

120 5.3. Global detectability of the programme

subtracted by 0.1 M!. As we can observe in Figure 5.5, there is no significant difference in the detection probability between these two approaches.

The errors in the radial velocities are composed of a photonic component, described by Equation 3.9, and a non-photonic component. The latter is mostly related to 1) the instrumental error, 2) other sources of noise intrinsic to the star (stellar activity, oscil- lation), and 3) perturbations caused by other non-detected companions. The impact of individual sources of non-photonic errors can not be easily evaluated separately, but we can estimate their global effect in the radial-velocity variations. To proceed this investi- gation, I selected a subsample of metal-rich and low-rotation stars for which no known planetary or stellar companion is present. For each star, I calculated the dispersion in the radial velocities and then I subtracted, quadratically, the mean photon-noise errors of the respective measurements. As a consequence, any residual variation in the radial velocities are probably due to the instrumental error, to other sources of noise intrinsic to the star, and/or to the presence of non-detected planets. Figure 5.3 shows the distribution of the radial-velocity dispersion corrected from the photon-noise errors. This distribution of non-photonic error sources is included, in the way described in the next subsection, in the simulations to estimate the detectability of the present programme.

It is interesting to notice that the photon-noise errors have a direct dependence on the stellar magnitude, spectral type, as well as the observing conditions (weather, air mass, seeing), since it is calculated based on the information of amplitude and width of the CCF profile, and also on the signal-to-noise ratio of the spectra. In such a way, the simulations reflect the real conditions of the observations. Nevertheless, in order to minimize the photon-noise errors, we normally chose the exposure time according to the weather conditions at the moment of the observation, also taking into account the stellar magnitude and spectral type.

5.3.1 Simulations: description

The grid of the simulations has 70 points in the orbital period P and 11 points in the planetary mass Mp, which results in 770 grid points. This is enough to allow a global view of the detection power of our programme. The orbital period ranges from 0.2 to − 3.3 in log P , which corresponds to about 0.6 to 1995 days. This is almost 3 times the duration of the radial-velocity measurements of the programme stars using the ELODIE spectrograph. For the planetary mass, the values are: M = 0.3, 0.5, 0.75, 1, 1.5, 2, 2.5, 3, 4, 5, 10 M p { } Jup

For each (P,Mp) grid point, I generated 1000 random Keplerian orbits, and for each orbit the following parameters were chosen as described:

Reference star: from the list of 221 metal-rich ([Fe/H] 0.1) and low-rotation (v sin i ≥

121 Chapter 5. Statistical study of the programme

Figure 5.3: Distribution of the radial-velocity dispersion corrected from the photon-noise errors for a list of 139 metal-rich and low-rotation stars for which no companion was detected. This distribution thus represents other sources of radial-velocity variation, like instrumental errors, stellar activity, or even non-detected planets.

6 km s−1) observed stars, I took only those for which at least 3 radial-velocity ≤ measurements were available. The resulting list contains 166 of such stars, which are selected randomly at the beginning of the simulations. For each selected star, the stellar mass M!, the Julian date of the measurements, and the respective photon- noise errors in the radial velocities are read.

Instrumental error and intrinsic stellar noise: from the list above of 166 stars, I se- lected only those for which no stellar or planetary companion was detected, resulting in a list of 139 stars. For each one of such stars, I calculated the radial-velocity dis- persion corrected from the mean photo-noise errors of the measurements. In the course of the simulations, a value of such non-photonic error sources is randomly selected and quadratically added to the photon-noise error of the chosen star.

Instant of periastron passage (T0): chosen randomly and uniformly in the range t1 and t1+P , where t1 represents the Julian date of the first radial-velocity measurement of the chosen star, and P is the orbital period corresponding to the selected grid point. In theory, this parameter can only be defined for non-circular orbits, but in the case of numerical simulations it is useful to allow a random change in phase of the Keplerian orbits created.

Orbital eccentricity (e): during the simulations of each grid point, I set the orbital eccentricity to be equal to zero, which is a reasonable consideration given the fact that the planets we are searching for (the hot Jupiters) have orbits in short periods, and by

122 5.3. Global detectability of the programme

consequence they mostly have low eccentricities. Nevertheless, I made a comparison of the results obtained by choosing different eccentricities: e = 0, e =0.5, and e =0.8 for values of Mp = 1 MJup and Mp = 5 MJup. The differences in the probability of detection is of the order of 10% or smaller between e = 0 and e =0.8, and smaller than 5% between e = 0 and e =0.5, for all the period range (see Figure 5.6). Orbits with high eccentricities are compensated by having steep radial-velocity variations, which are easier to detect. I noticed that, for non-zero eccentricities, the argument of periastron (ω) have to be defined. In such cases, values of ω are chosen randomly and uniformly from 0 to 2π.

Orbital inclination angle (i): values of sin i were chosen randomly and uniformly in the range 0 and 1.

With these parameters established, the semi-amplitude of the Keplerian orbit can then be calculated using the following relation:

1/3 2πG Mp sin i 1 −1 K1 = (m s ) (5.1) P (M + M )2/3 (1 e2)1/2 ! " ! p − where G is the universal gravitational constant. In a next step, simulated values of radial velocities are calculated at the same Julian dates in which the chosen star was observed taking into account the individual errors in the measurements. This is done by adding to 2 the theoretical radial velocities a random Gaussian of variance equal to ε(vr) , where ε(vr) is the error in each radial-velocity measurement.

Finally, the parameter E/I is calculated using the simulated values of radial velocities vr (simul) and their respective errors ε(vr). In this case, E is the dispersion of vr (simul) around its mean value and I is the mean value of the individual errors ε(v ). If E/I 2.0, r ≥ then I assume that the variation in the radial velocities is detectable. Otherwise, if the parameter E/I is smaller than 2.0, then I assume that any possible variation is not high enough to be detected so that the radial velocity is taken as constant.

The process just described is repeated for each one of the 1000 random orbits, and for each (P,Mp) grid point. Figure 5.4 shows the results by plotting the probability of detection of any important variation in the simulated radial-velocities (for which E/I ≥ 2.0) as a function of the orbital period of the planet, and for the different values of the planetary mass. It is interesting to notice in this figure that the probability of detection drops steeply for some values of the orbital period: P = 1 day is a typical interval between two measurements. P = 2 days is also a typical interval between two measurements, but it is specially related to our observation strategy, in which the first two measurements of each star were normally made separated by 2 days. We can also observe, even if less evident, a decrease for P = 30 and 60 days, which are related to the periodicity of the observing missions (and possibly to the Moon phases). Finally, the probability of detection is smaller around P = 365 days, which is a consequence of the Earth’s orbital movement as well as the result of seasonal changes in the observability of the stars.

123 Chapter 5. Statistical study of the programme

It is also clear in Figure 5.4 that the detection power of this programme decrease quickly for long periods, especially when P>100 days, limiting the detection of not very massive planets in long-period orbits. This is the result of the still small number of measurements of the target stars, and the relatively short duration of the programme if we compare, for example, to the CORALIE survey that has a long baseline of observations (almost 10 years) and for which most stars have at least 10 measurements. Nevertheless, we were mainly interested in the detection of planets in short-period orbits for which the probability of transit events is higher, and we can observe in the figure that we were able to detect at least 50% of the Jupiter-like planets in orbits shorter than 10 days.

At the beginning of this section, I noticed that we can not easily evaluate the impact of individual sources of non-photonic errors separately. Despite that, I tried to separate such error sources from the contribution of the instrumental error, which was estimated to be of about 6.5 m s−1 for the ELODIE spectrograph (see Naef 2004). Figure 5.7 compares the simulations when the uncertainties in the radial velocities are due to: 1) both photon-noise errors and non-photonic errors, the latter representing the contribution of the instrumental error, non-detected planets, and other sources of noise intrinsic to the star, like the stellar activity; 2) the photon-noise errors added quadratically to the ELODIE instrumental error of 6.5 m s−1. This figure gives us an idea of how the detection power of our programme is affected by the stellar activity and/or the presence of non-detected companions.

In Figure 5.8, I verify what the results would be if the number of iterations was 5000 instead of 1000, and if the grid of simulations had twice the number of points in the orbital period. I also plot in this figure the mean of 5 simulations with 1000 iterations each. We can observe that the difference between these simulations and those effectively used are not significant, which means that 1000 iterations with 70 points in the orbital period are sufficient to describe the global behaviour.

124 5.3. Global detectability of the programme

Figure 5.4: Probability of detection of radial-velocity variations as a function of the orbital period for several values of the planetary mass, from 0.3 to 10MJup. Notice the decrease in the probability for periods of 1, 2, and 365 days. A decrease for periods of 30 and 60 days, although less evident, is also present.

125 Chapter 5. Statistical study of the programme

Figure 5.5: Probability of detection of radial-velocity variations as a function of the orbital period comparing the results taking into account the stellar masses from Cox (2000) to those using stellar masses also from Cox (2000) but subtracted by 0.1 M! (see text for more details). I show the results at two regimes of planetary mass.

126 5.3. Global detectability of the programme

Figure 5.6: The same as Figure 5.5, but comparing the results of the simulations for three values of eccentricity at two regimes of planetary mass.

127 Chapter 5. Statistical study of the programme

Figure 5.7: The same as Figure 5.5, but comparing the results of the simulations by considering photonic and non-photonic (instrumental error, stellar activity, non-detected planets) error sources to those considering only photon-noise errors added quadratically to an instrumental error fixed to 6.5 m s−1 (see text).

128 5.3. Global detectability of the programme

Figure 5.8: The same as Figure 5.5, but comparing the results of the simulations for different values of the number of iterations and grid points: 1000 iterations with 70 points in the orbital period; the mean of 5 times 1000 iterations, also with 70 points; and 5000 iterations with 140 points in the orbital period. Notice that the results are very similar.

129 Chapter 5. Statistical study of the programme

Figure 5.9: Detection limits in the planetary mass at 50 and 90% for our ELODIE metallicity-biased search programme. The solid line is the limit computed assuming pho- tonic and non-photonic error sources, and the dotted line is the limit when only the photon-noise errors and an instrumental error of 0.0065 km s−1 are considered. Filled circles represent planets discovered or confirmed by ELODIE observations, whereas open circles indicate planets detected with HARPS or other instruments (see Table 4.2).

5.3.2 Detection limits in the planetary mass

We can use the same simulations previously described to estimate the minimum plan- etary mass for which a determined fraction of the random Keplerian orbits would be de- tected in the context of our programme. Such detection limits in the mass of the planetary companion are estimated for each orbital period in the selected grid of points. Figure 5.9 shows the results for the critical planetary mass above which 50% and 90% of the radial- velocity variations would be detected by our measurements collected with ELODIE during this programme.

130 5.3. Global detectability of the programme

5.3.3 Simulations: results

I used the simulations to determine the probability of detection of giant planets in the points (P,Mp) corresponding to each detection of stars with planets included in the present programme. I then used these probabilities to estimate the number of planets that could be effectively detected if any observational bias was not present. In order to allow a direct comparison with the estimates done in Section 5.1, I consider here only the stars with planets in our sample for which [Fe/H] 0.1 and K>30 m s−1. Table 5.1 ≥ lists the results for such stars and their planets. The parameter is the probability of Pdet detection of a planet given its mass and orbital period, and N (or 1/ ) is the number pl Pdet of planets (or a fraction of them) corrected from the observational bias that are expected to be detected in such a position in the (P,Mp) diagram. This time, I made 5000 iterations instead of 1000 in order to have a higher precision in the calculation of such probabilities. I estimated the errors in the number of detections by doing bootstrap simulations (Efron & Tibshirani 1993), which give results similar to those considering that the number of detection in a region of the (P,Mp) diagram follows the statistic of a Poisson distribution, namely σ(N ) N . pl ∼ pl We can notice- in Table 5.1 that the number of planets already found in this programme is not far from what we might expect if we only consider the planets having a probability of detection of at least 50%. The results of the simulations agree with the estimates based on the metallicity and period distributions of Sections 5.1 and 5.2: they all suggest that 7-8 stars with planets in the range of short periods, and at the same time likely to be detected with the present programme, are present in our sample of metal-rich stars.

Table 4.2 lists a total of 12 stars with planets. Two of them have planets with orbital periods longer than 100 days and, at the same time, are detectable with ELODIE. Many others are expected to be discovered at longer periods, but this will depend on the con- tinuation of our programme. In the next chapter, together with the conclusions of this work, I discuss the present situation of the programme, especially in which concerns the continuation of the observations using the new SOPHIE spectrograph.

131 Chapter 5. Statistical study of the programme

Table 5.1: Estimate of the number of detections (Npl) in the present programme corrected from the non-photonic error sources. The parameter is the probability of detection Pdet of a planet given its mass and orbital period (N =1/ ). The errors in the number of pl Pdet detection come from bootstrap simulations.

P<10 days:

Mp sin i P det Npl [MJup] [days] P HD 149143 b 1.36 4.09 0.66 1.52 HD 118203 b 2.13 6.13 0.74 1.35 HD 185269 b 1.03 6.84 0.50 2.00 subtotal 5 2 ∼ ± 10 P<100 days: ≤ HD 17156 b 3.12 21.20 0.68 1.47 HD 43691 b 2.49 36.96 0.54 1.85 HD 45652 b 0.60 43.70 0.08 12.5 subtotal 16 4 ∼ ± 100 P<1000 days: ≤ HD 75898 b 2.51 418.20 0.14 7.14 HD 132406 b 5.61 974.00 0.26 3.85 subtotal 11 3 ∼ ± Total 32 5 ∼ ±

132 Chapter 6

Conclusion and prospects

Contents

6.1 Summary of our results and main conclusions ...... 133 6.2 Newperspectives ...... 135

6.1 Summary of our results and main conclusions

In the present work we aimed to discover extrasolar planets based on a sample of 1061 solar-type and bright stars in the northern hemisphere. In order to increase our chances to detected such planets, we made use of a well-known phenomenon in which metal-rich stars are more likely to host giant planets in close-in orbits. Our strategy for the observations was to search for short-term variations in the radial velocities that could unveil the presence of a planet in a short-period orbit, the so-called hot Jupiters. The importance of hot Jupiters is due to their high probability to manifest a transit event, when the planet crosses the disk of its parent star as viewed from Earth. We have seen that the discovery and characterisation of transiting hot Jupiters orbiting bright stars allow us to derive some physical parameters of the planet, and of the system as a whole, that are not available only from radial-velocity measurements in a non-transiting system. Specially, the radius and mean density of the planet can be accurately determined, providing a more detailed study and description of its internal structure and chemical composition.

To date, 13 planets have been discovered orbiting stars present in our sample (see Table 4.2), and 2 of them were observed to transit the disk of their parent star. HD 189733 b (Bouchy et al. 2005) is a 1.15 Jupiter masses planet with a 2.2 days orbital period that was observed to transit the disk of its host star, a bright K dwarf, and for this reason

133 Chapter 6. Conclusion and prospects

it has became one of the most studied and well-known extrasolar planet. The transit of HD 17156 b in front of its companion has been recently announced (Barbieri et al. 2007), and will surely provide many interesting results in the near future. This planet is not a hot Jupiter, having an orbital period of 21 days, but the particular combination of its orbital parameters increased the probability of a transit event. Other hot Jupiters were unveiled among our target stars, but with no transit detection.

Seven planets orbiting stars in our programme have periods longer than 10 days. The discovery of extrasolar planets having long orbital periods may not provide the same diversity of physical information than that offered by transiting hot Jupiters, but will certainly contribute to statistical studies based on the global sample of known systems. The main discoveries related to our sample of 1061 stars is the following:

6 hot Jupiters (P<10 days): 3 discovered with ELODIE in the context of our main programme (HD 189733 b, HD 118203 b, and HD 185269 b), 1 only confirmed with ELODIE (HD 149143 b), 1 discovered with HARPS using a subsample of low decli- nation stars (HD 219828 b), and 1 confirmed with HARPS (HD 102195 b).

3 planets with 10 < P < 100 days: 1 discovered with ELODIE and also confirmed with SOPHIE (HD 43691 b), 1 discovered with ELODIE and confirmed with CORALIE (HD 45652 b), and 1 discovered by other groups of search for extrasolar planets and confirmed with SOPHIE (HD 17156 b).

4 planets with 100 < P < 1000 days: 1 discovered with ELODIE and also confirmed with SOPHIE (HD 132406 b), and 3 discovered by other groups (HD 75898 b, and HD 155358 b,c).

For some stars, we measured a huge radial-velocity variation that is probably related to the presence of a faint stellar companion, or just a brown dwarf star, or even a very massive planet. One group of such stars have also some measurements made with CORAVEL, the predecessor of ELODIE, and for 4 of them I was able to fit a Keplerian orbital solution to the available radial-velocity data (HD 62923, HD 89010, HD 212733, and HD 216320). The identification of possible candidates to brown dwarfs or very massive planets is of great significance as well, specially those having masses near the brown dwarf desert (the region between 20 and 60 MJup), since they improve the statistics and contribute to a better characterisation of the two populations.

The results of the simulations presented in Chapter 5 indicate that other planets are expected to be discovered in our sample: about 5 2 with periods shorter than 10 days, and ± 27 7 with longer periods. During the 2 years and 5 months of the observing programme ± using the ELODIE spectrograph, we were probably able to find most of the planets in short- period orbits, but the identification and characterisation of planets with longer periods requires more observations. Stars with no or a few number of radial-velocity measurements made during this time should also be observed or have the observations continued (at least

134 6.2. New perspectives

the metal-richer). The CORALIE spectrograph can be in charge of monitoring stars with small declinations, and such a subsample is defined in Chapter 3. Another subsample has been included in the planet search programme of the new SOPHIE spectrograph, the successor of ELODIE, and in next section I explain how this subsample was defined.

It is interesting to note that, according to the distribution of the number of planets as a function of the orbital period comparing CORALIE and the present ELODIE discoveries (see Figure 5.2), we found with ELODIE in less than 3 years of observations almost the same number of planets as that discovered by the CORALIE planet-search programme, but during about 10 years of radial-velocity measurements. This is a strong argument in favour of our observing strategy, in which only metal-rich stars were monitored: the high probability to find giant planets in close orbits around such stars allowed us to discover a large number of planets in a short interval of time.

6.2 New perspectives

The observation programme described in this work could not be finished since the ELODIE spectrograph was decommissioned before its conclusion. Almost 20% of the 1061 target stars has no radial-velocity measurements, and about 30% of the metal-rich stars has only 1 or 2 measurements. We then selected a subsample of stars that either have a small number of measurements, or show a significant variation in their radial velocities, to be observed using the new SOPHIE spectrograph. We established the following criteria to decide which stars should be included in the new subsample and which stars should not:

i) among the stars having at least one measurement made using ELODIE, so that a value for metallicity is available from the calibrations of the CCF profile, we selected those with [Fe/H] 0.1. Those having at least three measurements but with no ≥− significant variation in the radial velocities (E/I < 2.0) were not selected.

ii) for stars with no measurements, but with a photometric metallicity value available from Nordstr¨om et al. (2004), we selected those with [Fe/H] > 0.2. − iii) stars with declination δ 10 ◦ were already included in other subsamples to be ≤ observed either with CORALIE or HARPS, as described in Chapter 3, and were not selected to be monitored using SOPHIE.

iv) we did not select neither stars that revealed to be possible members of a binary sys- tem (those having a huge radial-velocity variation), nor those having a high rotation velocity (v sin i>6 km s−1).

In summary: among the initial sample of 1061 stars, 404 have a very low metallicity value, 93 was classified as constant in radial velocity, 98 are in the subsample of stars to

135 Chapter 6. Conclusion and prospects

be observed from the south hemisphere, and 33 are either binaries, or have a high rotation velocity. Thus, a total of 433 stars were selected to be observed using SOPHIE in the context of the French-Swiss consortium. As a matter of fact, some of such stars have been monitored with SOPHIE since the end of ELODIE, and in the present work I was able to show some results (see e.g. the case of the planets around HD 43691 and HD 132406, discussed in Section 4.4).

The higher precision in the radial-velocity measurements attained with the SOPHIE spectrograph (3-4 m s−1) in comparison with ELODIE (10-15 m s−1) will surely help us to complete the list of stars with planets that are likely present in our selected sample. Other improvements concerning the new instrument will also provide important consequences. In particular, a chromospheric activity study of stars observed using SOPHIE has been performed (Boisse et al. 2007), and will be soon implemented in the automated data reduction pipeline, helping us to elucidate the interpretation of the origin of the radial- velocity variation.

In complement to the new exoplanets that the SOPHIE and CORALIE observations may still provide based on the subsamples selected above, another approach can yield interesting results as well: a re-reduction of all available ELODIE data of our target stars using the same Data Reduction Software (DRS), accordingly adapted to ELODIE, as that used for the automated data reduction of the CORALIE, HARPS and SOPHIE observations. If a new treatment of the available spectra, using the new DRS, can improve the precision of the measurements made with ELODIE, then we may expect to identify additional stars having a significant and periodic radial-velocity variation, which will surely lead to the discovery of new extrasolar systems.

136 Part II

Appendices

137

Appendix A

Sample stars

139 Appendix A. Sample stars tive ec values ff N ) ! L (L 1.226 0 1.773 0 band, spectral V eff (K) T urements. i continued on next page sin v (km/s) ons are from Flower (1996). 0.060.47 1.4 1.3 5833 3.574 5665 1.003 1 1 0.11 5.10.10 6020 1.7 1.328 5708 1 1.239 1 2.53 4.8 6.409 1 0.48 1.7 5674 4.246 1 0.01 1.0 5623 0.608 1 0.03 2.2 5776 2.282 1 0.63 0.0 5785 1.226 1 0.61 4.8 5980 2.078 2 0.310.05 0.5 4.8 5717 0.867 5955 5.712 1 1 1.03 12.9 8.032 1 0.01 4.8 6134 1.918 1 0.050.44 1.1 3.2 5657 1.973 6136 1.576 1 1 0.370.07 3.0 2.0 5976 1.292 5445 0.527 1 1 0.54 16.30.11 6156 3.588 1.5 1 5823 2.091 1 0.34 1.0 5778 2.132 1 − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.083 0.083 0.275 0.03 0.0 5101 0.468 1 0.048 0.103 0.047 0.01 2.3 6073 1.402 1 0.067 0.080 0.051 0.130 0.096 0.051 0.165 0.28 5.2 5592 0.708 16 0.024 0.2850.098 0.07 0.28 1.5 3.1 5085 0.378 5894 1.955 3 3 0.080 0.21 2.1 5961 1.786 14 0.085 0.062 0.094 0.03 2.90.012 5811 2.868 19 0.037 0.109 0.14 2.6 5781 2.087 3 0.118 0.014 0.168 0.29 1.6 5581 1.703 3 0.052 0.25 3.90.007 0.090 6133 2.889 0.21 4 2.7 5904 0.863 3 0.038 0.165 0.081 0.032 0.072 0.259 0.06 0.1 5164 0.524 10 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 4.49 0.575 6 4.58 0.581 5 5.29 0.755 7 6.09 0.900 2 4.00 0.618 79 4.03 0.633 π 7.18 3.26 0.632 (mas) roper movement, apparent magnitude in the al. (2004). The stellar luminosity was calculated using the from the calibrations described in Section 3.3, and the e type Spectral V δ µ (as/yr) 75. Last column is the number of ELODIE radial-velocity meas 0.15975 7.490.16294 8.22 G0 G0 15.55 3.45 21.03 4.83 0.636 0.636 0.04321 8.48 K0 29.750.02217 8.18 5.85 0.888 G5 19.39 4.62 0.669 0.06350 7.63 G0 10.82 2.80 0.610 0.01672 8.47 G5 15.57 4.43 0.574 0.04321 8.36 G0 25.86 5.42 0.707 0.11077 8.19 F8 14.21 3.95 0.657 0.07924 8.10 G0 15.96 4.12 0.660 0.05827 7.84 F8 16.94 3.98 0.528 0.27639 6.970.00121 7.81 G0 F8 40.25 4.99 10.51 2.92 0.640 0.601 0.02382 8.16 G00.03810 8.360.02829 12.85 7.33 G0 3.70 F8 0.654 14.75 4.20 10.80 2.50 0.631 0.499 0.01789 7.110.17086 8.24 F8 G0 24.80 14.62 4.08 4.06 0.554 0.677 0.06131 7.950.07265 7.37 G5 F8 17.20 4.13 23.950.09467 7.19 4.27 0.691 0.504 G5 26.87 4.34 0.760 0.07034 7.56 F8 16.49 3.65 0.585 0.09602 7.86 F8 21.38 4.51 0.557 0.00566 8.060.01638 8.37 F8 F8 11.52 3.37 21.23 5.00 0.487 0.648 0.14838 7.67 G5 38.73 5.61 0.756 0.04309 8.29 F8 14.90 4.16 0.543 0.07219 8.02 K0 IV 34.52 5.71 0.870 . − − − − − − − − − − − − − − − − − − − − − − − − − − − ) are from the Hipparcos catalogue. The bolometric correcti =4 V α − ! µ bol (as/yr) B 0.00265 0.03412 0.03152 8.23 G0 18.60 4.58 0.581 0.00831 0.14089 0.03281 0.08994 0.01999 0.01806 0.01859 0.15976 0.16809 M − − − − − − − − − − − δ (d:m:s) , and adopting α BC (h:m:s) , and π , V HD 225097HD 70 00:03:25.9 +20:39:55.9 0.05724 00:05:41.6 +58:18:47.8 0.06598 HD 652 00:10:56.0 +48:06:37.5 0.16899 0.00093 8.41 G5 16.45 HD 804 00:12:28.3 +20:14:03.6 0.21641 Star ID HD 224983HD 225239 00:02:21.5 +11:00:22.4 HD 631 00:04:53.6 +34:39:35.3 00:10:39.1 0.76911 +12:49:12.9 0.10049 0.03898 6.09 G2V 2 HD 700 00:11:24.5 +23:49:05.6 HD 4899HD 5035 00:51:04.4 +02:44:36.8 00:52:40.1 +31:27:33.7 0.21327 HD 5515 00:57:18.2 +25:17:33.3 0.17686 0.01710 8.47 F8 16.6 HD 691HD 1497 00:11:22.4 +30:26:58.5 00:19:12.8 0.21073 +13:34:37.2 0.03526 0.01011 7.95 K0 V 29.3 HD 4635HD 4868 00:49:46.0 +70:26:58.2 00:51:29.1 +56:30:06.4 0.37041 0.06228 0.20147 7.75 K0 46.4 HD 3894 00:41:41.9 +20:08:13.5 0.08362 HD 5470 00:56:40.3 +17:57:35.5 0.09671 HD 4649 00:48:33.2 +00:22:55.4 0.06190 HD 3758 00:40:27.4 +16:34:47.0 0.05769 HD 1562HD 2435 00:20:00.4 +38:13:38.6 00:28:16.7 +28:03:37.8 HD 5651 00:58:26.4 +24:31:19.5 HD 2663 00:31:12.6HD 3161 +69:47:07.0 0.29768 00:38:14.7HD 4075 +84:40:29.9 00:45:03.9 +75:56:16.1 0.38921 HD 2330HD 2550 00:27:26.6 +26:16:43.7 00:29:23.6 +40:50:30.1 0.07067 HD 5312HD 5996 00:59:14.5 +83:06:09.7 01:02:57.0 +69:13:37.4 0.08203 0.22388 HD 3251HD 3628 00:35:52.3 +29:14:16.2 00:39:13.3 +03:08:02.1 0.10228 0.78134 0.29676 7.34 G2 V 21. HD 2520BD +320092 00:30:01.8 00:33:48.6 +75:14:34.4 +32:51:58.3 0.11504 0.00694 8.41 K0 13.1 HD 6047 01:02:56.3 +64:50:10.4 HD 3141 00:35:00.9 +42:41:41.2 Table A.1: The 1061 sample stars. The stellar coordinates, p type, parallax, and colour indexValues ( of metallicity andtemperature projected are from rotation theof velocity calibration come presented in Santos et

140 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 5817 01:03:39.4 +82:06:02.7 −0.17480 −0.06241 8.44 G5 22.56 5.21 0.600 −0.061 0.692 0 HD 6239 01:04:14.2 +55:50:48.0 0.12871 −0.06639 8.24 F8 13.66 3.92 0.489 −0.008 2.164 0 HD 6374 01:04:50.4 +07:02:45.4 0.15124 −0.10012 7.93 F8 13.07 3.51 0.543 −0.032 −0.14 3.4 6119 3.227 1 HD 6312 01:05:26.3 +63:43:16.7 0.22788 −0.16447 8.29 G0 18.19 4.59 0.600 −0.061 1.226 0 HD 6512 01:06:12.6 +13:15:09.9 0.16522 −0.01432 8.15 G0 20.84 4.74 0.656 −0.095 0.14 2.1 5849 1.102 3 HD 6660 01:07:37.9 +22:57:17.9 0.09908 −0.49171 8.41 K4 V 48.20 6.83 1.122 −0.519 −0.11 2.3 4456 0.237 3 HD 6697 01:07:51.5 +15:51:47.6 −0.08633 −0.03811 8.03 G5 III 16.81 4.16 0.725 −0.142 0.24 0.8 5669 1.962 3 HD 6664 01:07:58.8 +39:15:09.1 −0.07986 −0.00206 7.78 G1 V 26.79 4.92 0.621 −0.074 0.915 0 HD 6963 01:10:41.9 +42:55:54.6 −0.15350 −0.19995 7.66 K0 37.34 5.52 0.730 −0.146 −0.18 1.6 5481 0.563 1 HD 7355 01:13:51.6 +05:08:30.6 0.05985 −0.02159 7.69 F8 10.93 2.88 0.571 −0.045 0.39 4.2 6239 5.834 10 HD 7352 01:14:07.2 +25:49:39.6 0.04213 −0.06030 8.35 G0 V 14.33 4.13 0.550 −0.035 0.15 3.9 6214 1.828 27 HD 7510 01:15:41.2 +34:43:33.1 −0.05642 −0.02387 7.63 G5 16.95 3.78 0.781 −0.185 0.08 1.4 5431 2.897 2 HD 8100 01:21:04.0 +38:02:03.2 0.27769 0.01765 7.83 G0 20.25 4.36 0.661 −0.099 −0.04 2.2 5759 1.569 1 HD 7924 01:21:59.2 +76:42:37.0 −0.03412 −0.03445 7.17 K0 59.46 6.04 0.826 −0.221 −0.11 1.3 5220 0.374 1 HD 8249 01:21:59.7 +15:10:37.6 0.25642 0.01295 8.02 G1 III 17.12 4.19 0.667 −0.102 −0.09 1.9 5719 1.840 1 HD 8358 01:22:50.3 +00:42:43.4 −0.11411 −0.23936 8.26 G0 15.21 4.17 0.706 −0.129 −0.58 64.4 1.921 1 HD 8328 01:23:28.6 +46:53:26.3 0.10598 −0.04405 8.28 G5 12.69 3.80 0.691 −0.118 0.32 2.0 5810 2.674 2 HD 8365 01:24:03.5 +54:21:32.6 −0.09698 −0.02899 8.25 G5 17.06 4.41 0.590 −0.055 1.439 0 HD 8467 01:24:27.9 +39:03:43.4 0.21015 −0.03012 8.39 G5 31.77 5.90 0.784 −0.187 −0.11 0.0 5344 0.412 1 HD 8542 01:24:46.9 +20:09:31.8 0.11045 0.04273 8.15 G0 10.77 3.31 0.666 −0.102 −0.08 3.0 5726 4.138 1 HD 8553 01:24:53.9 +18:29:59.7 0.54540 −0.19137 8.49 K0 30.24 5.89 0.912 −0.297 −0.16 0.0 4957 0.460 6 141 HD 8745 01:26:53.8 +37:29:32.0 −0.01213 −0.08212 7.52 F8 18.57 3.86 0.549 −0.035 0.12 3.8 6205 2.344 4 HD 7585 01:29:01.2 +86:55:59.1 −0.01305 0.01065 8.40 G0 23.97 5.30 0.708 −0.130 0.679 0 HD 8730 01:29:01.8 +74:12:18.1 0.20118 −0.12483 7.24 G5 17.25 3.42 0.673 −0.106 0.04 2.5 5753 3.753 1 HD 9032 01:29:29.4 +29:32:36.2 −0.05159 −0.06197 8.41 F8 11.63 3.74 0.567 −0.043 −0.27 3.8 5982 2.638 1 HD 9081 01:29:42.4 +17:05:56.3 0.11702 −0.00496 8.13 G5 11.91 3.51 0.721 −0.140 0.26 2.4 5690 3.565 3 HD 9070 01:29:51.3 +31:00:26.4 0.22661 0.00099 7.93 G5 23.40 4.78 0.710 −0.132 0.35 1.5 5761 1.098 3 HD 9446 01:33:20.2 +29:15:54.5 0.19010 −0.05462 8.35 G5 V 18.92 4.73 0.680 −0.111 0.05 3.0 5734 1.128 1 HD 9536 01:33:43.5 +10:54:14.8 −0.02004 −0.11427 7.97 F8 18.99 4.36 0.527 −0.024 −0.47 4.2 6041 1.464 1 HD 9920 01:37:02.2 +05:37:50.8 0.12386 −0.07713 8.19 F8 10.41 3.28 0.618 −0.072 0.24 3.7 6017 4.138 2 HD 9776 01:38:04.7 +69:26:37.3 0.04948 −0.03162 8.41 G5 12.90 3.96 0.614 −0.069 −1.01 10.0 2.206 1 HD 9966 01:38:04.9 +30:47:23.5 0.28122 0.01355 8.18 F8 12.46 3.66 0.646 −0.089 −0.02 2.2 5816 2.962 1 HD 10183 01:39:44.0 +03:27:19.8 0.01121 0.00510 8.15 G0 10.66 3.29 0.548 −0.034 0.06 4.9 6184 3.959 1 HD 10211 01:41:00.6 +49:06:14.2 0.01728 −0.00025 7.40 F8 13.73 3.09 0.634 −0.081 4.970 0 HD 10337 01:41:06.6 +03:42:50.0 −0.10951 0.10706 8.09 G0 23.73 4.97 0.735 −0.149 0.12 0.0 5588 0.937 2 HD 10230 01:42:37.2 +70:10:25.7 −0.10085 0.06263 7.92 F8 10.68 3.06 0.506 −0.014 4.804 0 HD 10556 01:44:14.1 +54:53:11.1 −0.05014 −0.02097 7.27 F8 22.38 4.02 0.601 −0.062 2.074 0 HD 11087 01:49:10.8 +09:16:36.9 −0.02435 0.01648 8.18 F8 16.35 4.25 0.540 −0.030 −0.40 2.4 6023 1.629 1 HD 11045 01:49:16.1 +35:26:25.3 0.05212 −0.09858 8.14 G0 17.34 4.34 0.627 −0.077 −0.41 1.1 5719 1.566 1 HD 10790 01:49:56.4 +76:59:32.3 0.10113 0.03822 8.07 G5 25.18 5.08 0.856 −0.247 0.926 0 HD 11168 01:50:05.5 +14:21:04.9 0.00271 0.07296 8.50 G0 11.18 3.74 0.627 −0.077 0.08 3.0 5921 2.721 3 HD 11130 01:50:07.9 +29:27:52.4 −0.03583 −0.06384 8.06 K1 V 37.33 5.92 0.758 −0.167 −0.49 0.0 5267 0.397 1 HD 11271 01:50:52.1 +06:31:42.7 0.09545 −0.11398 8.45 G5 11.72 3.79 0.588 −0.054 0.21 3.3 6106 2.544 3 HD 11474 01:53:38.9 +39:52:59.2 0.06399 0.02863 7.73 F8 15.89 3.74 0.550 −0.035 −1.42 3.4 2.618 1 HD 11373 01:54:06.2 +66:10:34.3 −0.04493 −0.13921 8.47 G5 45.16 6.74 1.010 −0.392 −0.08 2.7 4734 0.230 2 HD 11616 01:54:11.9 +09:57:02.3 0.06901 −0.07054 7.78 G5 12.53 3.27 0.653 −0.093 0.00 3.3 5802 4.258 1 continued on next page Appendix A. Sample stars N ) ! L (L 0.7224.708 0 0 3.661 0 2.747 0 2.001 0 4.025 0 4.920 0 3.607 0 0.629 0 eff (K) T i continued on next page sin v (km/s) 0.07 4.3 6138 2.519 1 0.83 1.9 1.557 1 0.11 0.0 5649 1.112 1 0.36 5.7 6054 2.010 1 0.11 0.0 5091 0.429 1 0.06 2.9 6023 2.884 1 0.03 2.1 5198 0.386 1 0.17 4.6 6258 1.238 1 0.13 3.7 6043 1.743 1 0.00 1.2 5665 1.047 12 0.61 1.8 5873 2.800 1 0.46 2.0 6002 1.305 1 0.06 6.0 6121 2.325 1 0.00 4.5 6089 3.887 1 0.42 2.5 5882 4.567 2 0.22 2.4 6058 1.439 2 0.11 2.0 5353 0.447 1 0.25 0.5 5705 2.295 1 − − − − − − − − − − − − − − − − − − [Fe/H] 0.033 0.126 0.082 0.059 0.041 0.115 0.231 0.13 0.2 5284 0.587 3 0.087 0.04 4.5 5851 2.850 2 0.120 0.31 2.5 5793 2.148 4 0.028 0.260 0.050 0.237 0.102 0.06 1.8 5784 0.939 1 0.063 0.092 0.001 0.20 5.3 6523 2.491 2 0.012 0.057 0.11 4.9 6048 3.0390.040 8 0.043 0.091 0.05 1.50.120 5832 2.091 2 0.038 0.029 0.06 4.8 62230.029 2.441 13 0.095 0.29 4.4 5911 3.819 3 0.036 0.033 0.139 0.04 0.0 5602 0.7040.102 1 0.40 2.3 5923 1.632 3 0.133 0.07 5.4 5643 0.720 1 0.192 0.02 0.0 5380 2.4710.121 1 0.13 2.2 5715 0.921 5 0.044 0.074 0.04 7.3 5921 0.783 1 0.107 0.133 0.12 2.50.049 5664 1.842 14 0.100 0.06 3.4 5790 3.631 2 0.192 0.23 0.5 5466 0.700 5 0.035 0.128 0.16 1.9 5699 0.697 3 0.185 0.092 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 68 4.75 0.686 78 3.65 0.580 90 4.92 0.666 10 3.33 0.652 96 4.82 0.695 14 3.67 0.558 63 3.45 0.664 .83 4.04 0.650 π (mas) type Spectral V δ µ (as/yr) 0.02079 8.06 G0 13.90 3.78 0.546 0.00181 7.92 G0 11.13 3.15 0.635 0.06036 6.72 F8 21.65 3.40 0.596 0.44471 8.38 G5 15.38 4.31 0.564 0.15787 8.11 K0 V 30.87 5.56 0.838 0.03245 8.17 G5 12.78 3.70 0.643 0.00645 8.19 G0 14.66 4.02 0.536 0.04414 8.15 F8 12.58 3.65 0.461 0.003 0.20577 8.29 K0 35.21 6.02 0.845 0.13577 8.07 G0 15.77 4.06 0.604 0.03229 8.28 F8 12.45 3.76 0.470 0.04683 8.49 F8 16.15 4.53 0.501 0.01325 8.00 G0 13.18 3.60 0.593 0.06480 7.51 F8 12.88 3.06 0.562 0.12556 6.82 F8 29.74 4.19 0.566 0.04663 8.03 F8 14.300.10983 3.81 8.150.03991 0.537 7.28 G0 G0 18.50 4.49 16.71 3.39 0.539 0.656 0.08823 8.46 F8 12.06 3.87 0.552 0.03719 7.52 F8 14.95 3.39 0.545 0.03459 7.49 G0 23.21 4.32 0.666 0.01698 7.85 G5 30.00 5.24 0.711 0.06189 7.85 G5 16.66 3.960.07685 0.790 8.42 G3 V 20.28 4.96 0.696 0.06741 7.75 G2 V 29.40 5.09 0.622 0.02170 7.79 F8 12.78 3.32 0.568 0.31857 8.30 G8 V 25.82 5.36 0.674 0.04485 8.07 F80.04197 7.42 16.98 F8 4.22 0.711 14.00 3.15 0.578 0.22044 7.99 G5 29.35 5.33 0.790 0.23961 6.57 F8 36.57 4.39 0.551 0.13855 8.12 G8 V 26.89 5.27 0.705 0.06032 8.46 G5 29.52 5.81 0.781 0.21705 7.00 G3 V 24.38 3.94 0.651 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.20747 0.14448 7.99 G5 28.04 5.23 0.703 0.03221 0.01970 0.02643 8.41 G1 V 13.36 4.04 0.695 0.01608 0.12466 8.21 K2 V 35.00 5.93 0.871 0.00825 0.01302 0.02217 0.00949 0.03874 0.02168 0.07898 8.26 G0 25.23 5.27 0.720 0.08389 0.34290 0.11507 − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 19103 03:05:23.2 +32:05:48.0 0.03293 HD 18916 03:04:25.0 +50:01:56.6 HD 18768 03:02:52.2HD 18962 +47:06:38.0 03:06:49.4 0.06669 +67:24:37.8 HD 17820 02:51:58.3 +11:22:11.9 0.03733 HD 17225HD 18702 02:53:16.9 +78:40:14.6 03:00:36.1 +05:59:09.8 0.14401 0.69812 0.02498 8.17 G5 20. HD 17156 02:49:44.4 +71:45:11.6 0.09100 BD +600583 02:53:10.5 +61:05:49.1 HD 17499 02:48:46.8 +14:36:49.0 0.05787 HD 15670 02:32:22.1 +39:22:18.7 0.13882 0.08314 8.12 G0 12. HD 18450 02:58:52.4 +26:46:26.7 HD 17379 02:49:20.1 +53:01:23.0 0.04180 HD 15632 02:31:24.2 +22:05:14.9 0.14490 0.06220HD 16979 8.03 02:44:46.9 G0 +45:35:57.7 23. HD 15451 02:30:39.3 +43:40:44.4 0.01934 0.05221 7.74HD 17290 G0 02:49:02.1 13. +58:01:43.4 0.13545 HD 16663 02:41:46.3 +45:29:39.9 0.09170 HD 15819 02:32:49.5 +05:12:36.0 0.07922 HD 15397 02:30:14.0 +48:30:05.6 0.04084 HD 15069 02:28:20.6 +62:13:08.4 0.17164 0.04690 8.78HD 16067 G1 V 02:37:53.0 +65:19:07.2 17 HD 15292 02:33:40.2 +77:40:02.3 0.02043 0.12147 7.67 G5 26. HD 14305HD 14877 02:19:08.7 +19:41:15.7 02:24:39.8 +22:26:01.8 0.13508 0.00547 8.44 G0 11. HD 16086 02:35:10.5 +04:07:24.9 0.19647 HD 12382 02:01:27.9 +00:28:32.3 0.00612 HD 13908HD 14651 02:18:14.5 +65:35:39.7 02:22:00.9HD 15210 +04:44:48.3 0.02187 02:27:03.4 +06:13:35.3 HD 16175 02:37:01.9 +42:03:45.5 HD 15942 02:34:03.6 +12:10:51.1 0.18142 HIP 10679B 02:17:24.8 +28:44:30.4HD 15028 0.09815 02:25:47.1 +20:16:48.5 0.11086 HD 15851 02:35:35.5 +61:06:22.4 0.09768 HD 11850HD 12436 01:56:47.3 +23:03:04.1 02:02:39.3 +33:45:47.2 0.03782 continued from previous page Star ID HD 11638 01:54:24.6 +10:21:40.6HD 12741 0.10981 02:06:10.4 +46:51:35.6 0.01526 HD 13783 02:16:49.5 +64:57:09.4 HD 12536 02:03:01.7 +03:21:31.0 0.22025 0.13803 6.88 G0 20. HD 12800 02:09:07.9HD 13997 +71:33:07.2 0.30826 02:16:27.7 +12:22:47.2 0.22646 HD 13836 02:15:27.3 +27:21:26.2 0.28665 HD 13483 02:12:18.0 +21:58:57.7 HD 13403 02:12:56.0 +57:12:16.2 0.26714

142 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 19381 03:07:50.5 +25:33:07.0 −0.00831 −0.09952 8.28 F8 16.22 4.33 0.574 −0.047 0.01 4.6 6073 1.537 1 HD 19518 03:08:44.3 +15:20:05.4 0.06885 −0.19489 7.85 G8 V 24.44 4.79 0.642 −0.086 −0.13 1.5 5784 1.043 1 HD 19902 03:13:02.8 +32:53:47.3 0.18452 −0.06165 8.15 G5 23.73 5.03 0.732 −0.147 0.18 2.0 5622 0.885 3 HD 20065 03:14:04.8 +18:18:13.7 −0.01530 −0.22778 8.12 G5 24.52 5.07 0.624 −0.075 −0.34 1.6 5758 0.798 1 HD 20278 03:16:03.1 +11:37:42.4 0.19435 −0.01444 8.01 G5 10.74 3.17 0.608 −0.066 0.14 4.6 6009 4.554 3 HD 20678 03:20:59.4 +33:13:06.0 0.16655 −0.11228 7.94 K0 25.21 4.95 0.728 −0.145 0.11 3.4 5606 0.951 3 BD +180487 03:27:55.3 +18:52:56.4 0.01354 −0.06128 8.46 G0 14.32 4.24 0.568 −0.044 −0.29 1.1 5970 1.666 1 HD 21663 03:30:30.4 +20:06:11.7 0.15502 −0.06338 8.29 G5 21.61 4.96 0.739 −0.152 0.05 1.2 5547 0.948 1 BD +180492 03:30:56.1 +18:47:57.9 0.13132 −0.07085 8.12 F8 10.83 3.29 0.580 −0.050 −1.41 7.8 4.018 1 HD 21774 03:31:15.4 +20:46:05.1 −0.11236 −0.19642 8.08 G5 20.04 4.59 0.680 −0.111 0.30 1.5 5837 1.284 3 HD 21742 03:33:26.7 +59:25:00.3 0.16171 −0.30744 8.08 K1 IV 31.12 5.55 0.871 −0.260 0.14 0.8 5194 0.608 4 HD 21922 03:36:26.0 +67:57:07.8 0.06716 0.01707 8.20 F8 16.83 4.33 0.542 −0.031 1.515 0 HD 22481 03:37:20.7 +15:05:43.5 −0.02080 −0.06189 8.24 G0 14.67 4.07 0.534 −0.027 −0.35 2.7 6065 1.918 1 BD +630438 03:40:05.3 +63:52:28.6 0.13758 −0.13199 8.21 G5 23.92 5.10 0.818 −0.214 0.11 9.1 5334 0.882 3 HD 22718 03:41:25.7 +50:49:17.6 −0.03554 0.05960 8.20 G0 10.91 3.39 0.634 −0.081 3.771 0 HD 22872 03:42:41.6 +51:10:23.4 0.19906 −0.17833 7.95 F9 V 14.18 3.71 0.578 −0.049 2.726 0 HD 23221 03:43:24.7 +00:26:42.5 0.05006 −0.00233 8.48 G0 10.86 3.66 0.643 −0.087 0.02 3.8 5843 2.957 1 HD 23050 03:43:47.6 +42:36:12.1 0.36374 −0.25851 7.48 G2 V 23.46 4.33 0.583 −0.051 1.543 0 HD 23091 03:43:52.3 +37:58:41.5 −0.02585 −0.01798 8.23 F8 15.39 4.17 0.506 −0.014 −0.36 14.3 6161 1.728 1 HD 23074 03:45:21.1 +60:21:03.4 −0.05537 −0.00406 8.02 G0 12.87 3.57 0.581 −0.051 −0.07 5.1 6015 3.107 1 HD 23438 03:46:52.9 +41:51:33.8 −0.05822 0.04588 8.31 F8 15.32 4.24 0.606 −0.065 −0.41 1.2 5790 1.698 1 143 HD 23476 03:48:29.9 +57:46:03.4 0.07521 0.00161 8.50 G5 V 20.32 5.04 0.648 −0.090 0.832 0 HD 23565 03:48:37.5 +51:49:23.7 −0.07437 −0.04097 7.70 G5 V 16.24 3.75 0.656 −0.095 2.742 0 HD 24053 03:50:08.9 +06:37:14.4 0.05966 −0.04052 7.72 G0 30.74 5.16 0.684 −0.113 0.13 2.4 5754 0.761 12 HD 24301 03:52:59.4 +26:40:42.7 0.10694 −0.12303 8.00 G0 IV 11.45 3.29 0.644 −0.087 −0.24 2.0 5733 4.157 1 HD 24552 03:54:22.6 +01:15:41.9 0.02234 −0.11701 7.97 G0 22.13 4.69 0.631 −0.080 1.138 0 HD 24505 03:54:59.8 +28:11:17.2 0.01856 −0.05688 8.05 G5 III 13.78 3.75 0.737 −0.151 −0.07 0.4 5504 2.887 1 HD 24386 03:55:23.1 +53:33:27.3 −0.07579 0.02465 8.50 F8 V 12.77 4.03 0.534 −0.027 −0.37 3.5 6057 1.990 2 HD 24702 03:56:28.7 +22:40:27.9 0.17427 −0.23312 7.84 G0 21.41 4.49 0.688 −0.116 0.08 0.6 5721 1.414 1 HD 25100 04:01:06.0 +43:27:39.5 −0.06860 −0.03960 8.27 K0 28.69 5.56 0.810 −0.208 −0.04 0.0 5296 0.574 1 HD 25295 04:02:24.2 +31:19:50.4 −0.11339 −0.19252 8.32 G5 15.27 4.24 0.591 −0.056 −0.44 0.4 5829 1.684 1 HD 25461 04:03:50.8 +29:11:51.8 0.24771 −0.12833 8.18 K1 V 28.44 5.45 0.834 −0.228 −0.02 0.0 5234 0.647 1 HD 25682 04:04:30.7 +00:14:44.1 −0.20270 −0.35463 8.32 G5 21.76 5.01 0.769 −0.175 0.02 0.0 5443 0.925 1 HD 25869 04:06:13.0 +08:30:10.4 0.02339 0.10921 7.76 F8 15.89 3.77 0.462 0.002 2.462 0 HD 25979 04:07:08.7 +10:47:58.9 0.02856 −0.08194 7.97 F8 14.57 3.79 0.480 −0.004 2.430 0 HD 25173 04:07:23.7 +75:10:30.1 0.16784 −0.30780 7.22 F8 V 18.30 3.53 0.531 −0.025 3.148 0 HD 24894 04:08:04.3 +79:36:59.5 0.19937 −0.06341 8.29 F8 12.13 3.71 0.565 −0.042 2.709 0 HD 25890 04:08:27.6 +49:52:59.5 0.05924 −0.11254 8.06 F8 10.36 3.14 0.601 −0.062 4.664 0 HD 26421 04:12:23.0 +37:42:43.6 0.11745 0.09795 8.15 G0 12.75 3.68 0.617 −0.071 2.860 0 HD 26756 04:14:25.7 +14:37:30.1 0.10561 −0.01986 8.45 G5 V 21.91 5.15 0.693 −0.119 0.19 4.9 5750 0.772 3 HD 26018 04:15:13.6 +76:17:16.9 0.04256 −0.24115 8.19 G5 37.87 6.08 0.840 −0.233 −0.05 2.1 5204 0.364 1 HD 26900 04:16:32.7 +36:30:06.6 −0.02731 −0.22834 8.25 K2 39.08 6.21 0.918 −0.302 −0.04 2.7 4990 0.344 1 HD 27282 04:19:08.0 +17:31:29.1 0.11280 −0.02989 8.47 G8 V 21.24 5.11 0.721 −0.140 0.21 5.0 5669 0.817 3 HD 27495 04:22:00.3 +39:56:02.3 −0.02327 −0.11848 7.02 F8 19.51 3.47 0.616 −0.071 0.19 2.7 6003 3.471 3 HD 27685 04:22:44.8 +16:47:27.7 0.17330 0.00467 7.86 G4 V 26.96 5.01 0.677 −0.109 0.11 2.6 5768 0.870 2 HD 27741 04:23:43.9 +28:11:10.7 0.05659 −0.12700 8.31 G0 V 11.98 3.70 0.642 −0.086 −0.02 1.5 5830 2.847 1 continued on next page Appendix A. Sample stars N ) ! L (L 0.5931.886 0 0 0.764 0 1.862 0 1.397 0 1.689 0 0.394 0 1.413 0 2.434 0 2.176 0 4.365 0 0.685 0 0.615 0 3.266 0 2.968 0 1.095 0 5.445 0 eff (K) T i continued on next page sin v (km/s) 0.05 1.7 5837 1.604 1 0.27 4.6 6141 1.806 1 0.17 0.0 5611 0.799 1 0.26 0.0 5386 0.555 2 0.14 4.0 5915 3.442 1 0.18 0.0 5543 0.571 1 0.96 19.6 1.038 1 − − − − − − − [Fe/H] 0.083 0.132 0.022 0.115 0.13 2.4 5744 2.642 4 0.098 0.117 0.065 0.083 0.039 0.139 0.045 0.111 0.02 0.8 5722 0.788 1 0.035 0.13 2.9 6206 1.652 3 0.060 0.27 2.9 6097 2.858 2 0.161 0.231 0.14 0.2 5291 0.643 3 0.200 0.25 0.0 5444 0.780 4 0.060 0.07 3.7 6011 1.871 1 0.026 0.159 0.17 4.7 5572 0.580 25 0.060 0.17 3.6 6052 3.404 3 0.084 0.176 0.13 0.5 5485 0.943 3 0.106 0.03 2.3 5748 1.830 1 0.075 0.33 6.2 6037 1.932 11 0.049 0.19 5.2 6133 1.752 9 0.060 0.101 0.11 4.2 5807 1.108 11 0.110 0.109 0.28 2.1 5838 2.310 21 0.062 0.132 0.132 0.159 0.220.168 0.8 0.22 5592 1.7 0.786 5552 4 2.987 3 0.045 0.041 0.023 0.02 7.5 6253 2.019 1 0.079 0.161 0.030 0.21 3.9 6274 1.660 3 0.087 0.17 3.2 5904 4.044 3 0.040 0.100 0.22 3.4 5856 1.096 34 0.096 0.09 2.6 5825 1.113 1 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 39 4.06 0.464 0.001 92 5.45 0.710 76 5.14 0.660 71 4.42 0.571 24 4.13 0.599 87 3.81 0.533 14 5.50 0.747 11 5.27 0.679 84 5.41 0.710 97 3.51 0.571 69 3.61 0.563 .02 5.12 0.680 .55 5.49 0.710 π (mas) type Spectral V δ µ (as/yr) 0.19136 8.46 G5 14.85 4.32 0.636 0.13987 7.88 F8 17.75 4.13 0.522 0.03647 7.72 G0 16.54 3.81 0.687 0.26833 8.060.10972 G5 8.15 G0 25.65 5.11 15.76 0.690 4.14 0.607 0.40509 8.19 G5 34.88 5.90 0.720 0.07950 8.49 F8 V 14.10 4.24 0.550 0.14358 8.14 F8 12.76 3.67 0.598 0.20227 8.15 G8 V 30.22 5.55 0.750 0.29100 8.43 G5 25.50 5.46 0.837 0.31693 7.98 K0 28.08 5.22 0.800 0.01402 8.27 G0 11.00 3.48 0.599 0.03785 7.51 G0 19.75 3.99 0.639 0.09900 8.09 G5 23.94 4.99 0.770 0.20755 7.91 G0 18.15 4.20 0.673 0.04852 8.50 G0 13.26 4.11 0.623 0.02368 7.72 F8 19.68 4.19 0.578 0.04887 8.06 F8 IV 10.71 3.21 0.599 0.02472 8.30 G5 19.42 4.74 0.665 0.09215 7.12 G5 23.23 3.95 0.677 0.03745 8.44 F8 10.14 3.47 0.602 0.17417 8.340.22759 K0 7.40 G5 23.19 5.17 18.42 0.747 3.73 0.760 0.02915 8.45 F8 12.92 4.01 0.524 0.14156 7.83 G5 23.97 4.73 0.630 0.02163 8.21 G5 21.46 4.87 0.750 0.01996 7.78 F9 V 19.46 4.23 0.540 0.17741 7.66 G5 13.54 3.32 0.643 0.02095 7.67 F8 11.40 2.95 0.561 0.02782 8.10 G8 V 21.42 4.75 0.664 0.08964 8.05 G5 21.66 4.73 0.657 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.02457 0.03681 0.02224 7.82 G2 V 21.36 4.47 0.637 0.01350 0.00115 7.71 F8 20.08 4.22 0.560 0.02550 0.03325 − − − − − δ (d:m:s) α (h:m:s) HD 33313HD 33366 05:09:45.9 05:10:44.6 +07:10:53.7 +25:08:29.4 0.05164 0.01461 7.41 F8 21. HD 31865 05:03:33.3 +63:05:09.4 0.10112 0.31526 7.93 G0 31. HD 33588 05:11:31.0 +05:42:11.6 0.06065 HD 31864 05:03:30.7 +63:04:41.4HD 34031 0.10624 0.31091 05:15:11.6 7.63 +20:03:21.9 0.06467 G0 31. HD 32070 05:01:37.5HD 32016 +24:38:11.7 05:03:17.4 0.09373 +54:54:53.1 0.02325 HD 33866 05:13:25.1 +03:41:13.6 HD 32237 05:02:09.9 +14:04:53.7 0.08636 HD 30841 05:00:54.1 +77:46:31.8 HD 32274 05:01:57.2 +01:02:58.7 0.11500 0.04977 8.31 F8 16. HD 31867 05:00:17.5 +25:08:11.1 0.06205 0.00288 8.04 G2 V 26 HD 31781 04:59:38.8 +26:14:59.5 0.02014 HD 31219 04:58:11.5 +63:01:49.3 0.13528 HD 31501 04:57:59.3 +34:16:04.8 0.58116 HD 31452 04:56:10.7 +02:56:03.1 0.15061 HD 31338 04:55:41.7 +20:00:07.7 0.17392 HD 31354 04:55:22.7 +05:38:11.0 0.06270 0.03248 8.36 F8 14. HD 31016 04:54:33.4 +45:40:52.0 0.00233 0.01668 7.95 G0 14. HD 31000 04:53:56.2 +36:45:26.8 0.00617 0.01158 7.77 G5 35. HD 30925 04:53:07.0 +34:42:20.6 0.04349 HD 30974 04:52:45.1 +14:37:23.3 0.20005 HD 30376 04:51:57.0 +66:29:45.5 0.06000 HD 30518 04:51:27.1 +55:50:12.2 0.05275 HD 30572 04:49:48.0 +23:23:44.8 0.04484 HD 30589 04:49:32.2 +15:53:19.5 0.08737 HD 30467 04:49:05.5 +27:00:41.5 HD 30246 04:46:30.4 +15:28:19.4 0.09137 HD 30286 04:46:16.5 +03:16:07.6 0.05517 0.02983 7.81 G0 31. HD 29836 04:42:51.7 +18:43:13.6 0.10346 HD 29284 04:38:20.1 +37:21:44.7 0.00780 HD 29021 04:37:52.1HD 29400 +60:40:34.3 04:42:49.9 0.06062 +66:44:08.9 0.02206 0.35506 7.76 0.09150 G5 8.29 G8 V 33. 27 HD 276618 04:37:43.5HD 29862 +41:36:41.3 04:42:48.8 0.06371 +12:12:35.1 0.08019 HD 29356 04:37:26.7 +00:33:11.2 0.01723 0.01142 7.49 G5 15. HD 29355 04:37:26.1 +00:34:28.7 0.01521 0.01290 8.50 G5 16. HD 29037 04:35:35.4 +32:34:47.2 0.02895 HD 27757 04:32:15.5 +77:37:26.2 0.19927 HIP 20907B 04:28:52.6 +30:21:53.4 0.01418 HD 28635 04:31:29.4 +13:54:12.5 0.09260 HD 27969 04:26:49.4 +42:54:34.1 0.13475 HD 28097 04:28:54.3 +55:00:01.7 HD 28099 04:26:40.2 +16:44:48.9 0.11029 continued from previous page Star ID HD 27857 04:26:26.1 +52:30:17.2 0.07866

144 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 34720 05:20:03.6 +13:32:41.9 −0.02020 −0.08959 7.50 G2 III 18.93 3.89 0.561 −0.040 0.13 2.7 6167 2.291 2 HD 34887 05:23:03.1 +46:59:12.5 0.01274 −0.08471 8.15 G5 24.12 5.06 0.780 −0.184 0.28 0.7 5516 0.890 4 HD 35147 05:23:26.5 +16:58:53.4 0.00643 −0.15228 8.11 F8 19.97 4.61 0.606 −0.065 −0.47 2.3 5765 1.208 1 HD 35171 05:23:38.3 +17:19:26.9 0.25335 −0.00466 7.93 K2 69.76 7.15 1.104 −0.498 −0.28 3.1 4427 0.173 8 HD 35073 05:26:07.5 +61:14:29.4 0.09482 −0.25628 8.35 G5 20.64 4.92 0.700 −0.124 0.959 0 HD 35768 05:28:35.1 +32:51:57.7 0.05489 −0.04990 8.50 F8 11.58 3.82 0.556 −0.037 0.08 5.3 6164 2.437 3 HD 36248 05:30:47.4 +08:23:56.0 −0.01610 0.05569 8.05 F8 12.70 3.57 0.628 −0.078 0.14 3.8 5942 3.186 5 HD 36308 05:31:23.5 +12:33:21.9 0.17235 −0.13620 8.40 G5 28.00 5.64 0.814 −0.211 −0.10 3.0 5259 0.535 1 HD 35759 05:31:33.4 +64:19:07.6 −0.02841 −0.09190 7.74 G0 12.74 3.27 0.602 −0.062 4.138 0 HD 36548 05:33:09.1 +15:56:37.4 0.03164 −0.14083 8.19 F8 14.25 3.96 0.549 −0.035 −0.25 3.3 6053 2.138 1 HD 36387 05:33:16.9 +39:30:20.5 −0.06789 −0.13908 8.11 G0 18.52 4.45 0.617 −0.071 0.04 2.3 5938 1.407 1 HD 36130 05:37:33.5 +74:41:19.5 −0.02501 −0.18137 7.75 G0 19.96 4.25 0.617 −0.071 0.13 1.8 5975 1.692 3 HD 37008 05:38:12.1 +51:26:44.6 −0.54943 0.10581 7.74 K2 V 48.72 6.18 0.834 −0.228 −0.33 0.0 5106 0.331 1 HD 37216 05:39:52.4 +52:53:51.0 −0.01270 −0.14166 7.85 G5 35.91 5.63 0.764 −0.171 0.520 0 HD 37685 05:40:46.4 +09:15:55.6 −0.05323 −0.07003 7.96 G0 13.68 3.64 0.652 −0.092 −0.01 3.1 5801 3.026 1 HD 37773 05:41:34.3 +15:12:49.4 0.07464 −0.19960 7.67 G0 26.51 4.79 0.692 −0.118 0.02 0.6 5683 1.074 1 HD 37977A 05:44:24.7 +40:24:16.9 0.02152 −0.08480 8.28 G0 19.06 4.68 0.609 −0.066 0.10 5.0 5990 1.133 1 HD 37006 05:46:12.0 +78:15:22.5 −0.04511 0.06994 8.20 G0 28.20 5.45 0.722 −0.140 0.05 3.6 5600 0.597 1 HD 37393 05:46:29.6 +74:36:36.5 0.14251 −0.18713 7.33 G0 27.37 4.52 0.702 −0.126 1.388 0 HD 38400 05:47:49.9 +46:59:00.4 −0.11897 −0.09364 8.20 F8 13.60 3.87 0.593 −0.057 2.370 0 HD 247909 05:49:46.2 +32:58:23.4 0.01602 0.01215 8.46 K0 24.57 5.41 0.721 −0.140 0.10 1.4 5624 0.619 1 145 HD 39480 05:52:58.1 +03:02:52.9 0.03462 −0.09222 8.34 G5 12.21 3.77 0.655 −0.095 0.19 3.4 5873 2.692 2 HD 39392 05:53:19.0 +22:04:19.7 −0.07530 −0.07784 8.40 F8 11.25 3.66 0.552 −0.036 −0.53 2.0 5927 2.821 1 HD 39570 05:53:58.8 +12:25:09.9 −0.00844 −0.26015 7.76 F8 18.64 4.11 0.590 −0.055 0.01 2.3 6017 1.897 1 HD 40040 05:57:01.8 +15:44:29.0 0.08955 −0.24433 8.20 G0 14.69 4.04 0.654 −0.094 2.097 0 HD 40330 05:58:57.1 +18:58:47.5 0.04782 −0.04908 8.47 G0 13.50 4.12 0.606 −0.065 0.07 2.7 5987 1.897 6 HD 40647 06:06:05.8 +69:28:34.1 −0.12376 −0.05723 8.29 G5 31.30 5.77 0.783 −0.186 −0.04 1.9 5376 0.464 1 HD 40708 06:06:07.4 +67:38:23.9 0.03969 −0.31474 8.34 G5 27.00 5.50 0.771 −0.177 0.05 0.0 5449 0.590 1 HD 41708 06:07:59.0 +27:25:41.6 0.02589 −0.09432 8.02 G0 V 22.46 4.78 0.626 −0.077 0.05 1.5 5912 1.044 1 HD 41454 06:08:04.4 +50:34:59.2 0.04674 −0.07176 8.31 F8 12.99 3.88 0.461 0.003 2.222 0 HD 41750 06:08:24.7 +31:12:10.9 0.06840 −0.00865 8.28 F8 16.19 4.33 0.561 −0.040 −0.52 2.3 5900 1.528 1 HD 42012 06:09:56.4 +34:08:07.1 0.03354 −0.26018 8.44 K0 25.00 5.43 0.790 −0.192 −0.09 0.0 5334 0.638 1 HD 42160 06:10:01.0 +17:56:03.2 0.21011 −0.22899 8.48 G2 V 22.18 5.21 0.670 −0.104 0.720 0 HD 43062 06:14:24.4 +05:10:05.1 0.05804 −0.28551 8.39 K0 33.74 6.03 0.876 −0.264 −0.12 0.0 5073 0.392 1 HD 42820 06:14:54.7 +42:07:23.9 −0.01577 −0.07758 8.17 F8 V 15.67 4.15 0.589 −0.054 0.02 3.6 6025 1.826 1 HD 42819 06:16:12.4 +56:56:03.7 −0.20858 −0.18727 7.48 G0 18.90 3.86 0.599 −0.060 −0.02 1.8 5974 2.399 2 HD 43383 06:16:55.2 +25:29:44.2 −0.00797 0.02675 8.47 F8 V 12.68 3.99 0.540 −0.030 0.06 3.0 6212 2.070 7 HD 43691 06:19:34.7 +41:05:32.3 0.02423 −0.05314 8.03 G0 10.73 3.18 0.596 −0.059 0.22 4.7 6083 4.483 24 HD 43167 06:23:05.8 +74:50:45.9 0.05029 −0.02026 7.86 F8 12.34 3.32 0.524 −0.023 −0.03 3.1 6232 3.812 1 HD 45652 06:29:13.2 +10:56:02.1 0.20614 −0.06269 8.10 K5 27.67 5.31 0.846 −0.238 0.15 0.2 5269 0.743 19 HD 45205 06:30:15.6 +60:47:02.8 0.13663 −0.24710 8.45 G0 13.88 4.16 0.591 −0.056 1.813 0 HD 45743 06:32:20.0 +54:51:12.8 −0.04450 −0.11994 7.56 F8 10.59 2.68 0.545 −0.033 −0.06 4.6 6145 6.937 2 HD 45161 06:35:33.7 +76:51:58.9 −0.03824 −0.06126 8.14 G5 23.47 4.99 0.687 −0.115 0.891 0 HD 45821 06:36:30.6 +72:00:33.1 −0.04516 −0.24580 7.79 G0 20.03 4.30 0.700 −0.124 0.10 2.0 5690 1.697 2 HD 45875 06:38:43.1 +75:42:24.5 −0.03712 −0.10911 7.76 F8 19.28 4.19 0.505 −0.014 −0.36 2.7 6165 1.697 1 HD 48299 06:44:11.5 +34:09:59.9 0.03347 −0.07802 8.47 G0 13.26 4.08 0.567 −0.043 −0.06 3.5 6068 1.928 1 continued on next page Appendix A. Sample stars N ) ! L (L 2.164 0 0.535 0 3.107 0 0.926 0 5.465 0 1.419 0 1.617 0 2.439 0 eff (K) T i continued on next page sin v (km/s) 0.02 4.7 6243 2.132 2 0.13 3.9 5922 2.448 1 0.80 1.7 1.928 1 0.11 0.0 5305 0.460 1 0.03 0.4 5826 1.494 5 1.18 0.0 0.823 1 0.36 3.5 5976 1.304 1 0.260.51 8.6 2.8 6081 3.342 5908 4.203 1 2 0.40 6.0 6203 3.493 1 0.27 3.9 6066 4.455 1 0.33 0.0 4901 0.243 1 0.63 0.1 5663 1.420 1 0.03 0.0 5499 0.749 1 0.65 1.9 5956 1.374 1 0.09 2.7 5918 1.493 1 0.09 1.6 5837 1.038 1 0.37 1.6 5965 1.293 1 0.48 0.0 5691 0.665 1 0.41 0.0 5094 0.268 1 0.18 3.7 5807 5.143 1 0.62 3.1 5817 1.082 1 0.28 4.1 6006 1.379 1 0.10 4.9 6200 1.727 12 − − − − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.022 0.062 0.048 0.033 0.198 0.080 0.19 3.1 5949 1.459 3 0.086 0.029 0.180 0.23 1.1 5511 1.138 3 0.023 0.18 7.6 6318 2.613 3 0.161 0.038 0.161 0.07 0.1 5521 1.872 1 0.057 0.05 5.6 6020 1.595 2 0.001 0.030 0.039 0.063 0.03 2.2 5981 1.217 1 0.008 0.110 0.08 0.6 5750 0.802 1 0.054 0.077 0.292 0.032 0.341 0.03 1.2 4910 0.360 1 0.009 0.03 13.3 6369 3.308 1 0.071 0.156 0.025 0.090 0.065 0.107 0.02 2.7 5738 3.898 1 0.077 0.11 2.5 5936 2.243 2 0.080 0.077 0.039 0.109 0.06 1.1 5748 0.981 1 0.078 0.222 0.046 0.039 0.039 0.18 3.8 6191 4.442 3 0.032 0.066 0.09 2.6 5985 1.594 4 0.023 0.008 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 61 3.23 0.560 92 3.40 0.490 π (mas) type Spectral V δ µ (as/yr) 0.01426 8.11 G0 14.73 3.95 0.522 0.15528 7.93 F8 15.23 3.84 0.601 0.68388 6.74 F8 V 29.22 4.07 0.545 0.18721 7.97 G9 V 36.59 5.79 0.797 0.10661 8.00 G0 19.26 4.42 0.632 0.35276 7.37 G2 V 25.42 4.40 0.642 0.43667 8.48 F8 V 20.00 4.99 0.538 0.03869 8.15 G5 21.33 4.79 0.775 0.08871 8.49 F8 11.15 3.73 0.524 0.24697 8.24 G5 29.52 5.59 0.750 0.14143 7.62 G0 23.74 4.50 0.558 0.00767 8.31 K 15.29 4.23 0.750 0.25461 7.15 G0 26.89 4.30 0.594 0.02670 8.140.09418 8.14 G0 F8 11.93 3.52 11.63 3.47 0.471 0.540 0.05855 7.69 F8 24.06 4.60 0.603 0.21040 8.06 G0 25.63 5.10 0.679 0.08072 8.31 G0 20.88 4.91 0.627 0.04348 7.820.42558 8.19 F8 G5 11.70 39.91 3.16 6.20 0.543 0.959 0.29543 8.32 G5 44.92 6.58 0.907 0.07526 8.27 G0 17.17 4.44 0.617 0.05509 7.93 G5 12.77 3.46 0.493 0.06609 8.500.14303 G5 8.27 G5 22.07 5.22 17.30 0.744 4.46 0.649 0.21248 7.70 F8 22.23 4.43 0.530 0.10605 8.34 F8 16.11 4.38 0.607 0.09228 7.91 G2 IV 12.40 3.38 0.675 0.28786 8.00 G0 22.78 4.79 0.631 0.13543 8.15 G0 14.44 3.95 0.626 0.20654 7.80 G0 V 21.96 4.51 0.560 0.15781 8.02 G2 V 28.19 5.27 0.627 0.06229 8.07 G0 23.06 4.88 0.677 0.08888 8.40 K1 V 39.72 6.40 0.827 0.07452 7.84 F8 11.02 3.05 0.629 0.00571 8.01 G0 21.86 4.71 0.573 0.13897 7.77 F8 21.57 4.44 0.559 0.02135 8.37 F8 15.09 4.26 0.543 0.05769 8.48 G0 14.64 4.31 0.609 0.60361 5.44 F8 V 56.02 4.18 0.525 0.05815 8.40 F8 11.96 3.79 0.490 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.00729 0.04257 0.15251 7.08 F8 23.80 3.96 0.576 0.02611 0.03819 0.14847 0.04270 0.08725 0.14152 0.09288 0.02724 0.05605 0.07711 0.03836 7.65 F9 V 11.56 2.96 0.588 0.16456 0.00656 0.15721 0.01884 0.03298 0.05697 0.04358 0.02019 0.15851 0.05142 0.11200 0.40833 0.06176 0.04029 7.39 G0 14.35 3.17 0.560 0.11082 0.13305 0.00801 0.02029 0.09903 0.00385 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 61690 07:47:39.6 +73:59:30.5 HD 63107 07:47:32.6 +09:37:57.2 HD 60846 07:38:18.0 +42:27:35.1 HD 62301 07:44:56.2 +39:33:22.9 0.02778 HD 60272 07:36:58.1 +57:32:58.8 0.07261 HD 60521 07:36:14.1 +32:50:09.9 HD 60298 07:34:50.4 +24:57:16.1 0.12972 HD 59374 07:30:29.0 +18:57:40.6 0.02921 HD 59062 07:28:25.7 +04:09:11.5 HD 58900 07:28:07.3 +15:31:38.4 HD 51295 06:57:46.3 +22:53:33.2 HIP 34087B 07:04:05.8 +75:13:50.0 BD +481523 07:25:54.9 +48:30:44.3 0.07578 HD 50761HD 50476 06:58:14.3HD 51067A +58:36:07.0 06:59:07.8 07:04:04.0 +68:45:11.4 0.03383 +75:13:39.2 0.01945 0.01776 7.73 F8 12. HD 50483 06:53:54.1 +11:01:23.9 HD 53566 07:05:38.3 +03:56:58.0 0.03455 0.00770 7.84 F8 12. HD 57625 07:25:18.1 +56:34:09.0 HD 50060 06:51:55.9 +10:48:02.1 HD 51046HD 49985 06:58:37.8 +51:52:02.0 07:02:28.6 +79:09:51.9 HD 57901 07:23:47.0 +12:57:53.0 0.08459 HD 53505 07:06:50.2HD 53927 +32:47:22.1 07:08:04.2 0.20748 +29:50:04.2 HD 55647HD 56354 07:15:54.4 +43:33:42.7 07:19:20.8 +51:29:36.7 0.00004 HD 49838 06:51:21.9 +23:21:49.6 HD 54100 07:08:12.0 +15:31:15.0 HD 54684 07:15:47.9HD 56586 +70:30:21.2 0.01426 07:19:54.9 +46:58:20.0 HD 53075 07:06:53.4 +55:11:22.1 HD 49289 06:49:33.9 +40:10:31.9 0.06638 HD 56513 07:18:28.9 +27:15:10.2 0.15611 HD 54351 07:09:05.0 +15:25:17.7 HD 54046 07:08:00.3 +15:31:43.0 HD 54182 07:09:10.4 +30:23:13.0 0.02239 HD 49178 06:48:53.8 +37:30:18.6 HD 55458 07:13:53.2 +25:00:40.9 HD 54718 07:11:14.7 +32:36:54.0 HD 49385 06:48:11.5 +00:18:18.0 HD 54405 07:09:22.1 +16:32:58.2 HD 55088 07:11:43.8 +04:56:29.2 HD 54807 07:10:26.2 +01:38:34.9 HD 46588 06:46:14.4 +79:33:53.3 continued from previous page Star ID HD 47954 06:44:21.0 +55:52:04.4

146 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 62923 07:48:31.7 +47:45:53.7 0.05221 −0.23452 8.03 G5 19.43 4.47 0.740 −0.153 0.20 1.9 5606 1.490 15 HD 64324 07:54:48.5 +34:37:11.3 −0.12144 −0.17392 7.78 G0 28.33 5.04 0.659 −0.097 −0.00 2.2 5782 0.837 1 HD 64815 07:56:00.6 +06:30:44.0 0.20813 −0.09818 7.86 F8 12.64 3.37 0.626 −0.077 3.826 0 HD 65080 07:57:03.7 +01:30:47.6 −0.03597 0.02670 8.21 G0 18.57 4.55 0.638 −0.084 1.299 0 HD 65371 07:58:50.4 +10:07:47.1 0.44832 0.02380 8.12 K0 33.43 5.74 0.789 −0.191 −0.15 0.0 5313 0.479 1 HD 65523 07:59:35.7 +12:58:59.0 −0.23788 −0.04385 8.34 G5 29.46 5.69 0.770 −0.176 0.09 0.0 5469 0.495 1 HD 65629 07:59:53.4 +09:53:55.3 0.12368 −0.15136 7.98 G5 31.64 5.48 0.657 −0.096 0.558 0 HD 65857A 08:01:23.1 +16:56:58.0 −0.03100 −0.01582 8.42 G0 14.19 4.18 0.491 −0.008 1.703 0 HD 66485 08:04:42.8 +24:19:48.8 −0.03416 0.03702 8.43 G5 22.65 5.21 0.727 −0.144 −0.08 0.8 5531 0.747 2 HD 66550 08:04:52.2 +23:36:49.2 −0.12838 −0.08551 8.36 G5 26.60 5.48 0.806 −0.205 0.03 1.0 5336 0.617 1 HD 67150 08:07:26.0 +19:49:03.5 −0.09427 −0.04010 7.68 F8 10.34 2.75 0.599 −0.060 0.02 3.6 5991 6.668 1 HD 67346 08:08:08.7 +19:12:57.4 −0.03695 −0.09542 7.62 F8 14.49 3.43 0.619 −0.072 0.34 4.0 6054 3.604 3 HD 67481 08:10:44.8 +52:08:34.1 0.04775 −0.04646 7.41 F8 12.58 2.91 0.607 −0.065 −0.07 3.4 5926 5.781 1 HD 67517 08:11:01.4 +54:14:03.3 −0.00679 −0.06229 8.16 F8 V 10.31 3.23 0.602 −0.062 4.293 0 HD 68284 08:11:49.6 +04:16:28.3 −0.01014 0.05027 7.75 F8 13.47 3.40 0.589 −0.054 3.644 0 HD 68255A 08:12:12.7 +17:38:52.0 0.02829 −0.15094 4.67 G0 V 39.11 2.63 0.531 −0.025 0.22 4.5 6310 7.211 26 HD 67850 08:13:45.7 +63:06:14.3 0.14298 0.11340 8.42 G0 24.51 5.37 0.676 −0.108 0.624 0 HD 69056 08:15:33.2 +11:25:51.5 −0.19887 −0.23321 7.72 G5 26.45 4.83 0.731 −0.147 0.11 0.0 5597 1.064 3 HD 67869 08:17:43.3 +76:12:43.9 −0.06489 −0.11035 8.09 G0 10.49 3.19 0.623 −0.075 0.03 2.4 5914 4.508 1 HD 69475 08:18:46.2 +39:43:28.3 −0.00425 −0.12892 8.49 F8 13.05 4.07 0.518 −0.020 −0.36 3.2 6118 1.905 1 BD +012063 08:19:19.0 +01:20:19.9 −0.16459 −0.05301 8.35 G5 42.89 6.51 0.901 −0.286 −0.29 4.3 4934 0.257 1 147 BD +351801 08:19:27.8 +35:01:24.0 0.08938 −0.13825 8.30 G0 13.42 3.94 0.622 −0.074 0.09 3.5 5942 2.257 1 HD 68744 08:20:02.7 +73:20:10.1 −0.10039 −0.25065 8.44 G0 V 10.18 3.48 0.600 −0.061 3.407 0 HD 68788 08:20:05.2 +73:24:54.6 −0.31230 −0.48715 8.36 K1 V 32.22 5.90 0.853 −0.244 −0.06 0.0 5163 0.434 1 HD 69328 08:20:33.3 +65:23:38.3 0.01165 0.02155 8.43 G5 26.98 5.59 0.825 −0.221 0.00 3.4 5268 0.565 1 HD 70088 08:21:20.7 +34:18:35.9 −0.11645 −0.10992 8.48 G5 23.12 5.30 0.737 −0.151 −0.12 3.6 5483 0.692 1 HD 69960 08:22:16.1 +56:50:11.3 −0.02848 −0.01634 8.00 G5 16.34 4.07 0.756 −0.165 0.22 1.2 5565 2.178 2 HD 71053 08:25:55.0 +17:50:00.9 −0.08580 −0.02634 8.27 F9 V 12.75 3.80 0.596 −0.059 0.03 4.3 6005 2.533 1 HD 71227 08:28:09.0 +47:18:07.9 −0.01119 −0.04578 7.98 G0 V 22.29 4.72 0.639 −0.084 −0.08 1.4 5815 1.111 1 HD 71244 08:28:51.8 +53:51:44.0 0.08650 −0.05977 7.12 G0 23.73 4.00 0.565 −0.042 −0.40 4.3 5935 2.074 1 HD 71779 08:31:41.0 +53:33:38.0 −0.04962 −0.10145 8.13 G0 21.96 4.84 0.610 −0.067 0.08 2.5 5978 0.979 1 HD 72616 08:34:21.9 +16:52:48.8 −0.04928 −0.05794 8.35 G5 13.76 4.04 0.630 −0.079 0.36 2.0 6026 2.068 7 HD 72823 08:36:01.5 +30:01:04.8 0.02764 −0.00633 8.10 F8 11.73 3.45 0.533 −0.026 −1.26 9.7 3.392 1 BD +481666 08:36:39.5 +48:19:58.5 0.00275 0.01951 8.26 F8 15.47 4.21 0.489 −0.008 −0.31 15.9 6244 1.657 1 HD 71827 08:37:14.4 +77:02:48.4 0.03743 0.06827 7.29 F8 22.49 4.05 0.540 −0.030 −0.20 4.6 6105 1.959 2 BD +062008B 08:39:44.7 +05:46:14.0 0.16770 −0.30058 8.41 G5 26.98 5.57 0.810 −0.208 0.569 0 HD 73393 08:40:42.2 +55:40:03.9 −0.26456 −0.37002 8.00 G3 V 23.81 4.88 0.675 −0.107 0.08 0.0 5763 0.979 1 BD -002393 08:47:36.9 +00:01:04.4 0.15154 −0.09517 7.77 G0 17.66 4.00 0.655 −0.095 2.178 0 HD 74777 08:47:59.8 +50:16:38.5 −0.01021 −0.01967 8.48 G0 14.80 4.33 0.658 −0.097 0.15 2.8 5847 1.610 2 HD 75318 08:49:21.2 +03:41:02.4 −0.27385 0.00619 7.96 G5 30.16 5.36 0.746 −0.158 −0.15 0.0 5443 0.659 1 HD 75002 08:49:23.9 +50:40:25.9 −0.05557 0.00547 8.39 G0 11.28 3.65 0.576 −0.048 −0.02 2.8 6053 2.879 1 HD 75528 08:51:01.4 +15:21:02.3 −0.11813 0.07247 6.36 G2 IV 25.02 3.35 0.669 −0.103 0.09 3.1 5786 3.992 31 HD 75488 08:52:05.5 +47:33:49.7 0.04765 −0.20501 8.16 G2 V 17.21 4.34 0.557 −0.038 1.511 0 HD 75782 08:53:11.5 +37:04:13.0 −0.09741 −0.14484 7.09 G0 13.88 2.80 0.609 −0.066 0.30 4.0 6072 6.403 2 HD 75935 08:53:49.9 +26:54:47.6 0.01051 −0.00653 8.46 G8 V 24.66 5.42 0.766 −0.173 0.08 4.7 5477 0.633 1 HD 75898 08:53:50.8 +33:03:24.6 −0.09560 −0.02824 8.03 G0 12.41 3.50 0.626 −0.077 0.21 4.0 5977 3.395 22 continued on next page Appendix A. Sample stars N ) ! L (L 1.542 0 2.034 0 2.959 0 2.243 0 1.471 0 1.334 0 2.368 0 0.780 0 0.964 0 1.468 0 1.104 0 eff (K) T i continued on next page sin v (km/s) 0.34 4.0 6083 2.667 1 0.020.06 0.0 2.5 5725 1.610 5867 2.202 1 1 0.03 2.6 5611 1.614 1 0.09 2.0 5709 0.843 1 0.18 4.1 6068 2.769 1 0.25 0.0 56820.27 0.909 2.2 1 5889 1.441 1 0.03 0.0 57500.10 1.010 1.3 1 5545 0.596 1 0.17 10.1 4743 0.455 1 0.39 2.9 61630.03 2.057 2.6 2 5886 1.276 1 0.34 1.2 5675 2.148 1 0.19 2.2 5730 1.078 1 − − − − − − − − − − − − − − − [Fe/H] 0.012 0.14 5.0 6389 4.579 3 0.060 0.051 0.128 0.054 0.00 3.9 6020 2.430 1 0.025 0.037 0.107 0.077 0.130 0.049 0.067 0.10 2.4 5986 1.315 3 0.183 0.110.167 5.8 0.10 5450 1.7 0.450 16 5509 0.671 3 0.104 0.083 0.176 0.35 2.4 5575 0.806 3 0.036 0.057 0.172 0.30 1.7 5570 3.054 7 0.097 0.139 0.1040.033 0.13 0.32 2.3 8.2 5799 3.031 6298 23 2.865 2 0.101 0.094 0.03 2.8 5811 1.335 1 0.374 0.063 0.04 3.2 5982 3.142 1 0.2000.066 0.02 1.2 5350 0.540 2 0.090 0.07 4.5 5843 3.873 1 0.050 0.050 0.013 0.075 0.097 0.35 3.2 5926 3.970 6 0.090 0.144 0.19 3.6 5642 0.805 3 0.037 0.077 0.092 0.117 0.19 2.0 5759 1.798 3 0.177 0.06 3.6 5453 0.538 1 0.208 0.11 0.0 5357 0.569 2 0.068 0.19 3.4 6020 2.308 4 0.054 0.10 4.3 6058 3.003 3 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 09 4.34 0.598 28 3.37 0.649 π (mas) type Spectral V δ µ (as/yr) 0.05293 8.00 F8 10.50 3.11 0.500 0.00248 7.78 F8 17.78 4.03 0.583 0.06659 7.97 G0 13.97 3.70 0.705 0.04344 8.26 F8 13.08 3.84 0.588 0.05745 7.46 F9 V 17.80 3.71 0.530 0.07528 7.73 F8 17.25 3.91 0.554 0.31933 8.49 G5 14.79 4.34 0.674 0.08194 8.50 G0 14.85 4.36 0.708 0.04189 8.15 G0 17.60 4.38 0.578 0.11207 8.29 F8 17.64 4.52 0.610 0.24628 7.050.08878 K0 8.25 56.35 G5 5.80 26.29 0.779 5.35 0.758 0.04868 8.38 F8 16.92 4.52 0.636 0.26062 5.40 G8 IV-V 89.45 5.16 0.770 0.05081 7.90 F8 14.32 3.68 0.553 0.02179 7.690.15117 8.40 G5 G8 III 22.09 11.52 4.41 3.71 0.594 0.765 0.25236 8.36 G0 20.82 4.95 0.658 0.11714 8.480.02025 8.50 G5 G0 10.81 3.65 10.68 3.64 0.670 0.546 0.07452 8.30 G5 17.61 4.53 0.654 0.50762 7.200.11127 K3 V 7.69 57.05 G0 5.98 15.02 0.992 3.57 0.604 0.37255 8.140.08367 8.43 G5 G0 31.31 5.62 12.28 3.88 0.800 0.609 0.00152 7.84 F8 16.870.05273 8.17 3.98 0.502 G2 IV 10.88 3.35 0.659 0.14069 7.950.00662 G5 8.46 16.27 G0 4.01 21.56 0.649 5.13 0.727 0.11409 7.48 F8 23.87 4.37 0.556 0.12718 7.930.19414 7.86 G5 G0 23.23 4.76 18.78 4.23 0.651 0.690 0.01223 7.69 G5 38.21 5.60 0.771 0.10412 8.03 G5 32.18 5.57 0.810 0.20999 7.82 G5 16.50 3.91 0.611 0.04547 8.27 G0 11.68 3.61 0.589 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.14846 0.22908 0.093270.17883 0.04360 8.39 G0 13.07 3.97 0.626 0.09974 0.06469 0.10307 0.14751 0.020770.10133 0.00466 8.40 G0 21.31 5.04 0.670 0.08665 0.73005 0.07067 0.03721 0.14123 7.82 G5 33.54 5.45 0.720 0.18368 0.02463 8.34 G5 19.93 4.84 0.665 0.01805 0.06386 0.18726 0.02970 0.05219 0.07324 8.28 G0 V 22.77 5.07 0.580 0.05232 0.07025 8.25 G0 V 20.84 4.84 0.580 0.21474 0.00245 7.96 G0 20.85 4.56 0.624 0.07529 0.08916 0.05240 0.09426 0.02684 8.15 F8 20.65 4.72 0.626 0.14279 0.08531 0.05491 0.02350 0.11696 0.11044 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 84869 09:49:09.0 +36:31:47.4 0.01075 HD 84749 09:47:22.9 +02:18:54.7 0.07240 0.00851 7.72 G0 21. HD 84616 09:46:32.5 +01:43:08.2 HD 84209 09:43:43.8 +07:57:10.7 0.05045 HD 83921 09:42:31.3 +35:00:30.0 0.00046 HD 83186 09:40:35.1 +71:45:19.4 HD 83290HD 83408 09:37:25.2 +01:52:51.5 09:38:17.0 +01:50:06.3 HD 82606 09:33:05.6HD 82863 +04:10:17.1 0.04971 09:36:06.5 +47:31:24.6 HD 83147HD 82881 09:36:33.7 +07:07:21.6 09:37:34.6 +64:24:30.9 HD 82443 09:32:43.8 +26:59:18.7 HD 82460HD 82939 09:33:28.8 +46:13:43.2 09:36:04.3 +37:33:10.3 HD 82140 09:30:35.2HD 82885 +15:15:36.1 09:35:39.6 +35:48:36.5 HD 78437 09:19:58.6HD 80870 +82:36:45.6 HD 81505 09:23:10.4 +32:55:15.6 09:26:43.2 +26:21:02.5 0.15777 0.01453 HD 79726 09:16:07.7 +14:07:33.2 0.02719 HD 79282 09:14:14.6HD 80355 +33:49:00.8HD 80869 0.01049 09:21:03.0 +51:18:21.6 09:23:09.8 +33:54:19.2 0.07121 HD 78660 09:09:53.9 +14:27:24.3 HD 79969 09:17:53.4 +28:33:37.9HD 80914 0.04978 09:24:56.3 +59:59:35.2 HD 78233HD 79126 09:08:54.1 +50:47:05.8 09:12:42.5 +20:40:24.3 HD 80699 09:21:47.2 +18:44:46.8 HD 78536 09:08:53.9 +03:57:33.1 0.02073 0.00750 8.31 F8 10. BD +281698 09:08:27.2 +27:32:35.2 BD +281697 09:08:23.9 +27:32:07.6 HD 77599 09:05:46.0 +55:31:44.3 HD 76272 08:57:22.5 +54:50:46.3 HD 78277 09:08:02.7 +27:33:32.4 HD 76974 09:05:21.4 +74:41:48.2 0.14013 HD 76765 08:59:07.6 +31:54:57.4 HD 78317 09:07:33.1 +00:03:47.6 HD 73835 08:56:24.4 +83:53:14.6 HD 77006HD 76539 09:01:37.1 +49:44:12.6 09:03:30.1 +76:23:51.6 HD 76218 08:55:55.7 +36:11:46.3 HD 77278 09:01:47.6 +06:29:51.9 0.13292 HD 76025 08:54:18.7 +22:12:40.6 continued from previous page Star ID HD 75576 08:53:59.4 +62:35:01.6

148 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 85362 09:51:39.1 +17:28:32.8 0.00424 −0.18166 8.38 F8 14.96 4.25 0.631 −0.080 0.14 1.6 5932 1.706 3 HD 85238 09:51:55.4 +49:37:19.1 0.05102 −0.07915 7.80 G0 20.54 4.36 0.569 −0.044 0.10 2.4 6127 1.491 1 HD 85426 09:52:38.5 +35:06:42.0 0.26443 −0.14981 8.25 G0 16.35 4.32 0.663 −0.100 0.01 1.1 5773 1.629 1 HD 85689 09:53:48.4 +14:44:20.4 0.00095 0.04710 7.71 G0 22.13 4.43 0.583 −0.051 −0.29 2.8 5918 1.407 1 HD 85533 09:55:46.9 +70:02:28.1 −0.00414 0.02334 8.46 G5 18.34 4.78 0.717 −0.137 −0.03 0.5 5583 1.104 1 HD 85902 09:55:56.8 +38:04:54.1 −0.17069 −0.00049 8.12 G0 22.26 4.86 0.612 −0.068 −0.64 1.7 5675 0.962 1 HD 86133 09:57:02.3 +19:45:44.6 −0.23347 −0.00613 7.60 G0 23.54 4.46 0.560 −0.039 −0.17 1.6 6047 1.354 2 BD +202400 09:57:02.4 +19:45:14.1 −0.23274 −0.00783 8.40 G5 23.69 5.27 0.699 −0.123 0.694 0 HD 86165 09:58:15.8 +51:34:08.1 −0.16111 −0.05741 7.92 G0 18.99 4.31 0.625 −0.076 −0.23 1.3 5800 1.608 1 HD 86460 09:59:15.9 +27:31:22.9 −0.33037 −0.07805 7.78 G0 IV 24.37 4.71 0.582 −0.051 −0.21 1.2 5954 1.087 22 HD 86794 10:01:00.7 +08:56:56.4 −0.07956 −0.08695 8.06 G0 19.15 4.47 0.584 −0.052 1.358 0 HD 86661 10:01:38.3 +55:35:06.3 −0.17269 −0.44595 7.94 G8 IV-V 27.35 5.12 0.730 −0.146 −0.02 0.0 5546 0.814 1 HD 86884 10:03:12.0 +58:05:52.1 −0.17455 −0.13600 7.48 F8 13.58 3.14 0.583 −0.051 −0.27 4.2 5926 4.617 1 HD 87261 10:03:59.8 +07:10:34.1 0.05972 −0.02550 8.39 F8 10.20 3.43 0.594 −0.057 −0.30 3.5 5876 3.555 1 HD 86839 10:04:23.5 +70:52:00.6 −0.11034 −0.20717 8.41 F8 15.94 4.42 0.660 −0.098 −0.08 2.1 5746 1.483 1 HD 87680 10:07:14.8 +29:14:15.8 −0.00664 0.02683 7.96 G2 V 25.50 4.99 0.670 −0.104 −0.12 1.7 5697 0.882 3 HD 87666 10:07:21.6 +38:17:35.5 −0.16742 −0.07548 8.15 K0 20.83 4.74 0.664 −0.100 0.18 2.1 5839 1.107 2 HD 87912 10:09:46.0 +55:32:30.8 −0.22452 −0.12833 8.37 G5 12.53 3.86 0.727 −0.144 0.24 2.1 5663 2.592 3 HD 88232 10:10:39.8 +16:02:14.8 −0.22379 −0.10385 8.42 K2 22.46 5.18 0.760 −0.168 0.30 1.2 5585 0.786 4 HD 88446 10:12:19.1 +17:17:57.1 −0.15561 −0.22993 7.88 Gp 14.43 3.68 0.552 −0.036 −0.60 2.3 5899 2.769 2 HD 88402 10:12:56.0 +51:28:33.5 0.02088 −0.03468 7.57 G0 16.50 3.66 0.645 −0.088 0.13 3.1 5881 2.959 3 149 HD 88775 10:15:33.1 +47:22:10.0 0.01606 −0.10633 7.74 G0 13.10 3.33 0.695 −0.120 0.03 2.4 5678 4.130 1 HD 88909 10:16:14.3 +42:47:47.6 −0.00005 −0.06631 8.37 F8 15.15 4.27 0.566 −0.043 −0.25 3.1 5993 1.619 1 HD 89010 10:16:32.3 +23:30:11.1 −0.20023 0.03389 5.95 G2 IV 32.94 3.54 0.655 −0.095 0.11 2.6 5840 3.327 9 HD 89055 10:16:56.6 +25:51:38.6 0.16581 −0.29611 7.57 G0 V 28.11 4.81 0.595 −0.058 −0.04 0.5 5980 0.998 1 HD 89125 10:17:14.6 +23:06:22.3 −0.41329 −0.09722 5.81 F8 Vw 44.01 4.03 0.500 −0.012 −0.42 3.6 6158 1.962 2 HD 89110 10:17:54.6 +49:34:33.8 −0.06174 0.08672 7.69 F8 13.41 3.33 0.486 −0.006 3.719 0 HD 89652 10:21:41.9 +48:02:31.2 −0.21941 0.00673 7.59 G0 21.13 4.21 0.615 −0.070 −0.23 0.6 5834 1.754 1 HD 89813 10:22:09.5 +11:18:36.9 0.02132 −0.32428 7.78 G5 37.30 5.64 0.750 −0.161 0.04 0.0 5509 0.511 1 HD 90054 10:23:43.4 +02:37:09.6 −0.10513 −0.24964 7.87 G5 14.52 3.68 0.601 −0.062 0.38 4.3 6132 2.837 2 HD 90164 10:25:01.7 +30:22:16.5 0.09779 −0.12262 7.90 F8 V 18.55 4.24 0.567 −0.043 −0.20 3.0 6011 1.664 1 HD 90494 10:27:10.7 +19:48:48.2 −0.20593 −0.18150 8.39 F8 20.87 4.99 0.632 −0.080 −0.30 0.9 5748 0.863 1 HD 90508 10:28:03.8 +48:47:05.7 0.08056 −0.88176 6.42 G1 V 42.45 4.56 0.610 −0.067 −0.32 0.7 5814 1.267 1 HD 89881 10:28:26.3 +81:11:14.9 0.03486 0.01056 7.68 G0 16.80 3.81 0.662 −0.099 0.20 3.1 5854 2.604 2 HD 90681 10:28:51.4 +34:53:08.5 −0.11137 −0.06026 7.83 G0 22.00 4.54 0.652 −0.092 0.25 3.9 5908 1.321 6 HD 91148 10:31:45.6 +24:04:56.0 0.03303 −0.03898 7.93 G8 V 26.81 5.07 0.711 −0.133 0.11 1.8 5660 0.842 2 HD 91163 10:32:01.4 +29:43:57.5 0.08517 −0.01170 7.86 G2 V 10.65 3.00 0.614 −0.069 0.06 4.6 5956 5.341 10 HD 91332 10:33:09.7 +30:45:16.5 0.00749 −0.10065 8.25 G8 III 16.86 4.38 0.646 −0.089 0.32 3.9 5956 1.526 4 HD 91738 10:35:55.2 +18:02:44.0 0.06089 −0.05121 8.32 K0 26.10 5.40 0.790 −0.192 0.06 0.0 5396 0.656 1 HD 91702 10:35:56.1 +36:55:50.7 −0.06267 −0.12302 8.47 G5 19.21 4.89 0.724 −0.142 0.36 1.1 5721 1.002 2 HD 91988 10:37:36.8 +18:20:58.3 0.02935 −0.13505 8.38 G5 11.52 3.69 0.678 −0.110 0.13 2.7 5773 2.938 7 HD 92075 10:38:36.9 +40:19:19.7 −0.08518 −0.00539 8.40 F8 15.72 4.38 0.615 −0.070 −0.13 2.4 5875 1.500 1 HD 92242 10:39:17.5 +21:35:57.9 −0.13347 −0.14120 8.35 G5 14.84 4.21 0.619 −0.072 −0.16 1.9 5849 1.757 1 HD 92786 10:43:33.9 +48:12:50.9 −0.33083 0.18191 8.02 G5 39.56 6.01 0.749 −0.160 −0.34 0.0 5356 0.363 1 HD 92882 10:43:52.9 +27:24:07.6 −0.20854 −0.09773 8.14 F8 16.49 4.23 0.541 −0.030 −0.17 3.0 6114 1.660 1 HD 92903 10:44:12.2 +39:55:09.5 −0.14460 −0.00247 8.31 G0 17.49 4.52 0.591 −0.056 −0.24 1.4 5911 1.301 1 continued on next page Appendix A. Sample stars N ) ! L (L 0.416 0 1.086 0 2.662 0 1.138 0 0.487 0 eff (K) T i continued on next page sin v (km/s) 0.30 0.0 5568 0.776 1 0.57 2.9 6140 1.415 1 0.10 3.6 5969 1.970 1 0.16 5.2 5859 1.907 1 0.07 4.0 6085 2.208 1 0.02 0.0 5546 0.533 1 0.21 1.1 5785 0.904 1 0.05 1.5 5648 0.964 1 0.20 1.4 4860 0.254 1 0.01 4.2 6092 2.711 1 0.20 0.6 5742 0.903 1 0.37 3.2 5990 2.180 1 0.22 3.4 5954 4.059 1 0.00 2.4 6068 1.765 1 0.13 0.0 5580 0.750 1 0.73 1.3 1.108 1 0.73 1.6 1.845 1 0.53 2.0 5956 2.192 1 0.20 0.0 5365 0.626 1 0.37 2.7 5920 2.325 1 0.24 3.4 5743 5.181 1 − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.139 0.19 0.8 5664 1.583 2 0.104 0.25 2.6 5849 1.174 2 0.115 0.090 0.09 3.8 5852 5.395 1 0.319 0.007 0.056 0.071 0.040 0.146 0.120 0.26 3.1 5772 2.938 3 0.075 0.12 2.5 5947 1.387 14 0.063 0.18 6.0 6039 2.919 3 0.080 0.130 0.04 1.0 5643 0.946 1 0.090 0.120 0.083 0.02 1.9 5863 1.184 1 0.103 0.324 0.043 0.135 0.16 0.6 5671 0.611 5 0.182 0.12 0.6 5460 0.739 3 0.089 0.087 0.25 4.0 5934 3.522 8 0.036 0.043 0.06 4.1 6117 1.982 2 0.051 0.077 0.12 5.1 5940 4.643 7 0.047 0.128 0.103 0.08 0.3 5785 2.405 1 0.051 0.229 0.11 1.6 5284 0.420 2 0.060 0.105 0.07 3.1 5768 4.111 1 0.035 0.098 0.07 0.8 5807 2.329 1 0.039 0.12 4.2 6167 3.278 28 0.032 0.172 0.078 0.02 2.4 5889 2.826 1 0.085 0.08 2.5 5877 2.818 2 0.229 0.046 0.086 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M 72 4.94 0.632 π (mas) type Spectral V δ µ (as/yr) 0.01061 7.98 K0 19.12 4.39 0.720 0.00813 8.05 G5 V 21.19 4.68 0.670 0.15113 7.95 K0 27.36 5.14 0.687 0.01016 7.80 G0 11.00 3.01 0.649 0.20917 7.85 F8 20.20 4.38 0.488 0.05217 8.39 F8 13.66 4.07 0.591 0.06567 8.35 G0 14.23 4.12 0.616 0.00595 7.55 F8 V 18.90 3.93 0.561 0.13283 8.45 K0 26.70 5.58 0.730 0.10696 8.37 G0 11.62 3.70 0.695 0.21442 7.94 G0 20.22 4.47 0.624 0.09246 8.43 G0 11.07 3.65 0.604 0.05189 8.16 G0 22.69 4.94 0.707 0.20403 8.21 G5 20.30 4.75 0.648 0.02895 8.17 G5 22.27 4.91 0.694 0.04571 8.39 G0 12.01 3.79 0.668 0.06308 8.48 G0 11.14 3.71 0.566 0.15402 8.37 K0 25.70 5.42 0.714 0.24243 8.26 G5 21.80 4.95 0.646 0.23788 7.19 F8 22.40 3.94 0.553 0.07056 8.47 G0 13.08 4.05 0.567 0.03980 8.24 G0 10.20 3.28 0.581 0.11512 7.96 F8 10.99 3.16 0.626 0.04526 8.46 G0 13.90 4.18 0.574 0.51546 8.27 G5 24.23 5.19 0.705 0.12716 8.40 G0 12.59 3.90 0.668 0.09045 8.34 G5 32.79 5.92 0.835 0.06632 7.93 G5 22.29 4.67 0.599 0.07898 7.97 G5 11.77 3.32 0.672 0.21646 7.29 F8 23.27 4.12 0.549 0.08185 8.50 G5 24.27 5.43 0.765 0.03504 8.40 G0 V 11.51 3.71 0.640 0.11144 8.07 K0 34.51 5.76 0.835 0.00837 7.99 F8 15.06 3.88 0.573 0.10278 7.90 G5 10.71 3.05 0.641 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.00808 0.21347 0.05567 0.19091 0.03114 8.31 G5 34.77 6.02 0.937 0.02813 0.00789 0.18787 0.45104 0.15135 0.06485 0.15216 0.20770 0.30082 0.05287 8.25 G5 19.03 4.65 0.637 0.21400 0.21329 0.04065 8.35 K0 43.91 6.56 0.942 0.15579 0.15912 0.27742 0.03993 7.69 G5 32.73 5.26 0.777 0.11750 0.06940 0.02661 7.35 G0 16.74 3.47 0.644 0.13055 0.06293 0.20649 0.12414 0.06349 0.13008 7.86 F8 23.26 4.69 0.581 0.20283 0.10850 0.05828 0.05030 0.10495 0.07761 8.05 G5 15.00 3.93 0.660 0.08663 0.06355 8.39 G0 V 10.53 3.50 0.560 0.21843 0.08493 8.33 F8 13.17 3.93 0.544 0.11691 0.00106 8.21 G0 12.52 3.70 0.629 0.06274 0.19032 0.08554 0.06813 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) continued from previous page Star ID HD 93153 10:45:29.3 +01:00:38.8 HD 93215 10:46:09.2 +25:45:39.7 0.05317 HD 93330 10:46:33.3 +04:38:44.1 HD 93664 10:49:25.9 +41:23:24.1 HD 93811 10:50:40.4 +51:47:59.0 HD 94012 10:51:13.8 +09:13:28.1 HD 94291 10:53:13.4 +16:58:31.7 HD 94383 10:54:10.6 +34:07:52.8 HD 94426 10:54:25.5 +30:31:55.1 0.06046 HD 94718 10:56:16.4 +27:44:22.5 HD 94667 10:57:09.2 +65:16:41.1 0.00098 HD 94835 10:57:09.6 +21:48:17.5 HD 94861 10:58:15.4 +61:57:01.7 HD 94880 10:58:21.5 +59:16:53.4 0.01260 HD 95177 10:59:31.6 +07:13:12.1 HD 95072 10:59:06.2 +40:59:01.0 0.21725 0.02593 8.26 G5 21. HD 95516 11:01:41.4 +20:49:54.2 HD 95980 11:04:18.9 +05:47:44.5 HD 95544 11:05:19.7 +81:02:21.1 HD 96612 11:08:14.1 +38:25:35.9 HD 96691 11:08:26.4 +21:04:52.7 HD 96853 11:09:37.9 +41:54:49.4 HD 96937 11:09:40.2 +02:27:22.5 HD 97167 11:11:37.6 +42:49:03.4 HD 97140 11:11:43.6 +58:53:59.2 HD 97194 11:11:48.9 +42:49:55.6 HD 97178 11:11:51.6 +52:05:51.6 0.07240 HD 97286 11:13:06.4 +68:46:04.7 0.00212 HD 97657 11:14:45.3 +43:19:45.5 0.04939 HD 97890 11:15:51.3 +34:36:26.3 HD 98078 11:17:05.3 +21:17:47.1 HD 98248 11:18:13.9 +15:58:10.1 HD 99233 11:25:31.8 +42:37:58.4 HD 99303 11:25:40.0 +20:00:07.7 HD 99419 11:26:27.2 +20:31:05.2 HD 100069 11:31:00.7 +09:37:12.5 HD 100446 11:34:07.1 +65:14:33.1 HD 100667 11:35:25.3 +53:34:26.8 HD 100796 11:36:05.5 +30:43:17.1 HD 101305 11:39:28.5 +02:50:47.5 HD 101444 11:40:30.9 +03:39:14.8 0.15332 HD 101904 11:44:00.3 +64:57:51.1 HD 102161 11:45:31.5 +25:06:39.7 HD 102195 11:45:42.3 +02:49:17.3 HD 102618 11:48:51.3 +18:49:23.1 HD 102732 11:49:44.6 +06:31:23.6

150 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 103072 11:52:08.3 +18:45:18.7 0.03137 −0.30054 8.40 K2 39.48 6.38 0.860 −0.250 −0.20 0.0 5085 0.281 1 HD 103150 11:52:45.6 +47:32:47.0 −0.04736 −0.09038 8.44 F8 10.36 3.52 0.536 −0.028 3.186 0 HD 103126 11:53:41.7 +86:13:49.3 −0.22272 0.24125 8.28 G5 32.39 5.83 0.720 −0.139 −0.33 0.0 5450 0.420 1 HD 103431 11:54:35.1 +19:25:40.2 −0.45268 −0.01539 8.43 G5 27.44 5.62 0.760 −0.168 −0.22 0.0 5372 0.524 1 HD 103828 11:57:22.6 +58:17:59.3 −0.10213 −0.02895 8.44 G0 18.47 4.77 0.664 −0.100 −0.04 1.3 5749 1.076 1 HD 103847 11:57:29.0 +19:59:02.1 −0.38504 0.05344 8.07 G5 34.39 5.75 0.830 −0.225 0.09 2.7 5291 0.490 1 HD 103913 11:58:04.3 +25:08:16.2 0.07371 −0.00258 8.27 F8 12.15 3.69 0.572 −0.046 −0.05 3.2 6055 2.769 1 HD 104243 12:00:14.4 +05:21:48.7 −0.29700 −0.12859 8.40 G5 27.81 5.62 0.759 −0.167 0.04 1.6 5481 0.523 1 HD 104289 12:00:41.2 +59:21:11.3 0.02516 0.00971 8.06 F8 15.90 4.07 0.529 −0.025 0.07 4.9 6255 1.914 7 HD 104389 12:01:16.2 +79:05:14.8 −0.06782 −0.01238 8.45 G0 16.80 4.58 0.607 −0.065 −0.09 2.5 5918 1.242 1 HD 105279 12:07:16.0 +10:31:53.0 −0.03607 −0.07415 8.40 G5 12.49 3.88 0.661 −0.099 0.27 3.7 5886 2.441 3 HD 105320 12:07:37.8 +32:01:41.6 0.03618 0.01772 8.30 F8 V 14.58 4.12 0.526 −0.023 −0.44 2.7 6057 1.825 1 HD 105421A 12:08:07.2 +55:27:50.7 −0.17806 −0.01689 7.79 F8 19.47 4.24 0.513 −0.018 −0.15 4.8 6222 1.626 1 HD 105422B 12:08:09.8 +55:27:53.5 −0.18035 −0.03253 8.45 F8 21.54 5.12 0.567 −0.043 −0.12 3.7 6043 0.740 1 HD 105585 12:09:16.2 +42:28:54.1 −0.13320 −0.01085 8.48 G5 IV 18.43 4.81 0.664 −0.100 0.29 3.7 5885 1.038 3 HD 105844 12:10:59.7 +22:51:42.5 0.20224 −0.07486 8.07 G5 23.31 4.91 0.773 −0.178 0.20 0.0 5505 1.017 5 HIP 59431B 12:11:26.9 +53:25:07.4 −0.17598 −0.13761 8.14 K0 31.70 5.65 0.880 −0.268 −0.08 4.1 5078 0.559 2 HD 106366 12:14:08.4 +30:49:09.1 −0.13107 −0.10590 8.25 G0 15.23 4.16 0.640 −0.085 −0.05 1.1 5824 1.862 1 HD 106423A 12:14:27.3 +08:46:56.6 0.05476 −0.13119 7.50 G0 19.08 3.90 0.616 −0.071 0.40 3.8 6089 2.336 2 HD 106510 12:15:01.5 +30:08:29.5 −0.00658 0.01943 8.27 F8 V 19.65 4.74 0.582 −0.051 −0.55 1.4 5815 1.058 1 HD 106679 12:16:07.2 +03:33:35.9 −0.07749 −0.01472 8.04 F8 13.32 3.66 0.454 0.005 −0.23 5.7 6406 2.716 1 151 HD 106811 12:16:38.5 +63:37:46.8 −0.27706 0.06233 8.32 F8 14.30 4.10 0.625 −0.076 0.10 2.5 5936 1.952 1 HD 106888 12:17:36.1 +14:26:34.2 −0.10228 −0.03637 8.17 F8 17.10 4.33 0.543 −0.032 0.14 5.8 6234 1.516 2 HD 107146 12:19:06.5 +16:32:53.9 −0.17565 −0.14828 7.04 G2 V 35.07 4.76 0.604 −0.063 0.04 4.5 5982 1.050 1 HD 107211 12:19:24.8 +39:37:28.9 −0.14391 0.01329 8.38 G2 IV 14.85 4.24 0.665 −0.101 0.39 2.5 5922 1.755 2 HD 107303 12:20:06.1 +15:42:53.3 −0.07042 −0.06183 8.04 F8 15.42 3.98 0.499 −0.012 −0.24 4.4 6236 2.055 1 HD 107582 12:21:28.3 +61:44:50.1 −0.30060 −0.26221 8.23 G2 V 24.27 5.16 0.622 −0.074 −0.69 0.0 5621 0.734 1 HD 107728 12:22:41.0 +05:54:53.8 0.00530 −0.00788 8.47 G0 18.82 4.84 0.617 −0.071 −0.07 4.6 5893 0.983 1 HD 108024 12:24:39.4 +05:02:05.6 −0.03346 −0.26547 7.97 F8 20.97 4.58 0.578 −0.049 −0.37 1.3 5903 1.223 1 HD 108076 12:24:46.0 +38:19:07.5 −0.58737 0.06453 8.03 G0 V 26.94 5.18 0.585 −0.052 −1.06 0.2 0.706 3 HD 108436 12:26:53.8 +69:43:46.2 −0.10746 −0.18033 8.47 G0 18.83 4.84 0.658 −0.097 −0.43 0.0 5608 1.006 1 HD 108575B 12:28:04.8 +44:47:30.5 −0.17889 −0.00223 8.12 K 23.61 4.99 0.680 −0.111 0.16 5.6 5779 0.888 24 HD 108713 12:29:15.0 +34:42:12.4 −0.11065 0.02512 7.57 F9 V 10.80 2.74 0.574 −0.047 −0.06 4.1 6044 6.650 1 HD 108793 12:29:33.1 +56:19:55.0 −0.00526 −0.18538 7.73 G5 16.82 3.86 0.635 −0.082 0.04 1.7 5878 2.448 1 HD 108942 12:30:42.0 +50:58:28.3 −0.22427 −0.18874 7.91 G5 21.02 4.52 0.697 −0.122 0.20 2.1 5741 1.383 4 HD 108984 12:31:18.3 +20:13:04.1 −0.01747 −0.16960 7.91 K0 39.62 5.90 0.863 −0.253 0.05 0.8 5180 0.438 1 HD 109202 12:32:30.5 +56:58:34.8 −0.01358 −0.12151 8.12 G0 15.35 4.05 0.633 −0.081 0.17 2.6 5938 2.053 2 HD 109303 12:33:19.0 +49:18:06.8 −0.09827 −0.04732 8.14 F8 12.03 3.54 0.517 −0.019 −0.66 3.3 5998 3.102 1 HD 109413 12:33:34.1 +72:28:44.8 0.09030 −0.03635 7.50 F8 20.83 4.09 0.571 −0.045 −0.09 5.6 6042 1.914 1 HD 109552 12:35:28.0 +28:50:14.9 0.06078 −0.06266 8.09 F8 IV 10.76 3.25 0.605 −0.064 0.19 3.7 6040 4.223 3 HD 109625 12:35:57.9 +31:43:57.5 −0.02664 −0.00622 7.95 G0 15.67 3.93 0.553 −0.036 −0.73 3.9 2.200 1 HD 110313 12:40:03.9 +68:48:08.8 −0.43845 0.03061 7.88 F8 22.25 4.62 0.610 −0.067 0.05 2.2 5966 1.199 1 HD 110408 12:41:16.4 +58:57:43.2 −0.06789 −0.05486 8.08 F8 12.42 3.55 0.533 −0.026 0.09 5.6 6249 3.093 1 HD 110463 12:41:44.5 +55:43:28.9 0.12153 −0.00436 8.27 K3 V 43.06 6.44 0.955 −0.337 −0.07 2.4 4880 0.288 1 HD 110514 12:42:38.9 +02:34:36.1 −0.01034 −0.05065 8.04 K0 32.71 5.61 0.795 −0.196 0.543 0 HD 110869 12:44:29.1 +58:41:27.3 −0.21311 0.02799 8.01 G5 V 22.10 4.73 0.662 −0.099 0.13 1.4 5825 1.116 3 continued on next page Appendix A. Sample stars N ) ! L (L 2.512 0 0.982 0 1.279 0 0.856 0 1.059 0 0.880 0 2.613 0 2.012 0 0.843 0 1.280 0 1.026 0 1.002 0 3.260 0 1.324 0 2.669 0 0.214 0 eff (K) T i continued on next page sin v (km/s) 0.51 1.8 5918 1.036 13 0.25 2.2 5945 1.500 1 0.11 0.0 5639 0.652 1 0.20 1.1 5782 0.788 1 0.18 0.6 5770 0.828 1 0.02 1.1 5870 1.085 1 0.02 0.0 5454 0.590 1 0.03 2.7 5829 1.427 1 0.63 1.1 5754 1.968 1 0.19 2.2 5918 1.340 1 0.00 0.3 5254 0.431 1 0.86 1.1 1.043 1 − − − − − − − − − − − − [Fe/H] 0.038 0.048 0.15 3.6 6127 2.245 3 0.126 0.15 2.6 5701 0.799 5 0.132 0.24 2.1 5716 3.506 2 0.050 0.080 0.050 0.075 0.19 1.8 5979 2.218 3 0.037 0.061 0.071 0.11 2.6 5967 4.492 2 0.133 0.27 1.6 5722 1.397 3 0.080 0.16 1.9 5937 1.318 9 0.047 0.01 3.5 6073 1.990 1 0.116 0.081 0.071 0.07 4.2 5954 1.661 3 0.102 0.071 0.085 0.137 0.02 1.3 5604 0.885 1 0.079 0.023 0.079 0.105 0.113 0.29 4.5 5820 2.919 4 0.082 0.30 3.3 5984 2.943 4 0.048 0.118 0.214 0.090.124 0.0 0.37 5328 1.8 0.645 5801 1 2.430 3 0.112 0.18 3.0 5781 1.790 2 0.168 0.086 0.052 0.043 0.055 0.119 0.17 1.3 5741 3.594 2 0.123 0.15 4.7 5714 4.219 56 0.104 0.07 1.7 5775 1.519 1 0.015 0.058 0.026 0.225 0.036 0.438 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .93 3.52 0.710 .26 4.84 0.692 .62 3.91 0.700 .67 4.74 0.552 π (mas) type Spectral V δ µ (as/yr) 0.10606 6.29 G7 V 58.23 5.12 0.703 0.06852 7.56 F8 17.71 3.80 0.579 0.11860 8.39 F8 V 15.60 4.36 0.580 0.10274 7.66 G0 IV 18.22 3.96 0.623 0.06622 8.48 G0 19.94 4.98 0.600 0.07837 8.07 F9 V 10.56 3.19 0.617 0.02614 8.14 G8 III 18.92 4.52 0.712 0.15204 8.46 G0 16.37 4.53 0.632 0.06197 8.15 G0 15.12 4.05 0.574 0.01364 8.13 G5 16.90 4.27 0.616 0.12154 8.42 G5 18.80 4.79 0.667 0.17292 8.29 G2 V 22.39 5.04 0.640 0.36489 8.24 G8 V 22.69 5.02 0.717 0.20593 8.42 G5 18.34 4.74 0.630 0.08992 7.67 F8 16.33 3.73 0.526 0.04346 7.69 G5 29.57 5.04 0.671 0.03935 8.40 G0 16.84 4.53 0.576 0.08826 8.12 K2 III 29.08 5.44 0.817 0.09132 8.34 G5 III 15.10 4.23 0.682 0.02896 8.35 G5 III 16.56 4.45 0.641 0.28827 7.96 G0 16.67 4.07 0.590 0.13279 7.81 G5 III 13.60 3.48 0.693 0.07843 8.05 K0 11.29 3.31 0.699 0.13024 8.19 G5 17.43 4.40 0.670 0.11685 8.41 F8 16.20 4.46 0.507 0.02799 8.47 G5 15.98 4.49 0.595 0.06862 7.65 F8 16.30 3.71 0.533 0.09553 8.47 K2 47.63 6.86 1.051 − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.35975 0.13932 5.95 G0 V 57.57 4.75 0.557 0.12381 0.00660 8.18 G2 V 14.03 3.92 0.575 0.33455 0.14455 0.03208 8.48 G0 18.78 4.85 0.632 0.23215 0.06631 0.07357 8.13 G0 18.99 4.52 0.554 0.24522 0.02978 0.00303 0.26082 0.14828 8.15 G5 V 27.30 5.33 0.689 0.03662 0.01170 8.50 G0 20.75 5.09 0.634 0.19157 0.36289 0.13591 0.02707 8.39 G5 20.58 4.96 0.616 0.16778 0.05278 8.44 G5 13.37 4.07 0.630 0.26538 0.07495 0.02148 8.34 G5 11.78 3.70 0.684 0.04096 0.06346 7.96 F8 13.83 3.66 0.635 0.42506 0.28370 0.037780.12045 8.48 G5 25.29 5.49 0.760 0.13159 0.04499 0.01860 8.07 G2 V 22.19 4.80 0.585 0.01733 0.02422 8.01 F8 12.59 3.51 0.566 0.08784 0.06301 0.02065 0.09129 0.01799 0.07079 7.77 K0 IV-V 42.12 5.89 0.830 0.06737 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) continued from previous page Star ID HD 110897 12:44:59.5 +39:16:44.1 HD 110950 12:45:25.3 +29:46:33.6 HD 111395 12:48:47.1 +24:50:24.8 HD 111470 12:49:27.1 +05:10:19.7 0.11788 HD 111938 12:52:18.6 +46:46:40.0 HD 111814 12:51:45.6HD 112299 +19:13:43.8 12:55:28.2 0.01857 +25:44:17.1 0.01240 0.03629 8.14 G5 11 HD 112001 12:53:01.6 +26:47:39.2 HD 112756 12:58:55.2 +05:06:23.2 HD 113893 13:06:35.8 +11:02:40.3 HD 113621 13:04:19.0 +50:04:05.0 0.07595 HD 114036 13:07:15.5 +37:15:56.9 HD 113998 13:07:24.9 +11:08:05.5 0.05174 HD 114000 13:07:30.4 +03:56:13.5 HD 114060 13:07:39.7 +24:00:35.3 HD 114216 13:07:53.4 +62:58:59.3 HD 114285 13:09:15.2 +23:53:34.9 HD 115231 13:15:37.0 +09:00:57.7 HD 115273 13:15:59.2 +04:38:48.5 HD 115349 13:16:11.3 +35:53:09.1 0.26761 HD 115339 13:16:19.9 +27:44:03.8 0.07257 HD 115382 13:16:48.3 +12:24:55.6 0.07873 HD 115519 13:17:20.0 +32:29:15.6 0.12242 HD 115692 13:18:41.8 +07:11:11.6 HD 115755 13:19:06.7 +02:54:12.4 HD 115954 13:19:56.6 +38:22:08.6 HD 116091 13:20:07.8 +59:51:43.1 HD 116057 13:20:38.8 +36:40:32.5 0.10172 HD 116056 13:20:26.9 +43:06:40.2 HD 116272 13:21:12.2HD 117243 +61:05:05.3 13:28:39.3 +28:26:54.9 HD 117122 13:27:27.5 +44:42:25.4 0.08168 0.06313 8.42 G5 19 HD 117655 13:24:38.7 +83:18:30.6 0.02458 0.09258 7.24 G5 21 HD 117576 13:30:31.0 +44:34:58.2 HD 117845 13:31:33.8 +58:57:10.5 HD 117697 13:31:58.2 +08:58:37.5 HD 117858 13:32:30.9 +36:02:06.5 0.08332 HD 118051 13:33:49.8 +32:20:13.0 0.00415 HD 118203 13:34:02.5 +53:43:42.7 HD 118936 13:36:40.4 +75:54:31.3 0.25886 HD 118687 13:37:34.9 +41:26:05.4 HD 118790 13:38:10.6 +45:23:44.7 HD 119056 13:40:40.7 +19:12:00.5 HD 119332 13:41:13.4 +56:43:37.8 HD 119802 13:45:14.8 +08:50:09.6 HD 120162 13:45:19.2 +68:59:58.1 0.11617 0.05203 7.96 F8 22

152 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 119824 13:45:22.9 +11:28:45.4 0.03999 −0.00779 8.30 G0 22.14 5.03 0.657 −0.096 0.844 0 HD 120065 13:46:47.7 +22:31:25.9 −0.17918 0.05276 8.29 F8 12.54 3.78 0.584 −0.052 2.563 0 HD 120553 13:49:51.2 +14:53:02.1 −0.21298 −0.14809 8.15 G5 22.86 4.95 0.594 −0.057 0.877 0 HD 120877 13:51:57.2 +14:01:25.6 0.01832 −0.10035 8.14 F8 15.03 4.02 0.560 −0.039 2.030 0 HD 121131 13:53:05.2 +27:48:24.4 0.23996 −0.39528 8.37 K1 V 34.99 6.09 0.815 −0.212 −0.30 0.0 5174 0.354 1 HD 121234 13:54:09.9 +03:09:14.2 0.01408 −0.06368 8.21 G0 13.33 3.83 0.543 −0.032 2.403 0 HD 121320 13:54:28.2 +20:38:30.4 0.21041 −0.07599 7.89 G0 30.20 5.29 0.687 −0.115 0.676 0 HD 121763 13:56:10.6 +48:21:41.3 −0.16649 −0.03072 8.37 K0 17.14 4.54 0.713 −0.134 0.26 1.5 5715 1.373 3 HD 121979 13:56:17.9 +66:56:41.0 −0.13867 0.05499 8.45 K0 22.52 5.21 0.770 −0.176 0.24 3.5 5530 0.770 25 HD 121825 13:56:55.0 +44:16:57.1 −0.02363 0.00818 7.65 G0 22.47 4.41 0.612 −0.068 −0.28 1.2 5823 1.456 1 HD 122727 13:59:50.8 +73:23:53.3 −0.00810 −0.01770 7.80 G5 12.90 3.35 0.687 −0.115 0.26 3.9 5798 4.036 19 HD 122237 14:00:03.6 +27:08:27.5 −0.03329 0.05966 8.34 F8 13.17 3.94 0.630 −0.079 2.268 0 HD 122444 14:01:52.1 +00:25:20.3 −0.08983 −0.04384 8.32 F8 15.60 4.29 0.540 −0.030 1.570 0 HD 122518 14:01:57.5 +21:33:22.0 −0.12106 0.01065 7.93 G0 18.15 4.22 0.614 −0.069 1.736 0 HD 122813 14:03:06.1 +40:06:25.1 −0.07354 0.04407 8.33 G0 11.00 3.54 0.563 −0.041 3.165 0 HD 123300 14:04:51.7 +58:52:03.1 −0.17882 0.01941 7.58 G5 20.94 4.18 0.635 −0.082 0.10 3.2 5902 1.823 3 HD 123710 14:04:57.3 +74:34:24.9 −0.14722 0.09012 8.21 G5 22.84 5.00 0.590 −0.055 0.836 0 HD 124330 14:11:17.3 +54:24:32.9 −0.14718 −0.03694 7.88 G4 IV 17.36 4.08 0.652 −0.092 0.28 1.7 5920 2.017 3 HD 124244 14:12:17.1 +08:24:25.1 0.11782 −0.09566 8.47 G5 10.46 3.57 0.651 −0.092 0.09 1.8 5845 3.227 2 HD 124441 14:12:59.3 +31:34:39.5 −0.03385 −0.02769 8.23 G0 III 13.89 3.94 0.511 −0.017 −0.06 5.0 6267 2.142 1 HD 125193 14:15:59.9 +56:41:20.9 −0.00299 −0.09135 6.68 F8 21.01 3.29 0.580 −0.050 0.28 3.4 6163 4.018 17 153 HD 125141 14:16:05.4 +50:10:09.3 −0.30028 0.18607 7.78 G5 22.15 4.51 0.615 −0.070 0.06 1.7 5953 1.330 1 HD 125272 14:16:26.2 +58:04:59.3 −0.14654 0.12258 8.07 F9 V 16.13 4.11 0.594 −0.057 −0.06 2.0 5975 1.900 1 HD 125039 14:16:29.0 +26:37:44.1 0.02889 0.05622 8.36 F8 12.87 3.91 0.465 0.001 2.166 0 HD 125347 14:16:29.5 +63:03:51.5 0.03382 0.03312 8.32 F8 12.19 3.75 0.456 0.004 −0.55 6.1 6267 2.503 1 BD +640992 14:17:14.4 +63:48:25.5 0.03945 0.01032 8.35 F8 11.98 3.74 0.510 −0.016 −0.17 4.9 6225 2.573 1 HD 125231 14:18:06.5 +01:08:45.7 0.03785 0.05540 8.44 G0 13.86 4.15 0.537 −0.029 1.785 0 HD 125320 14:18:08.1 +26:47:43.5 0.00560 −0.05765 8.20 G5 IV 13.00 3.77 0.766 −0.173 2.892 0 HD 125505 14:19:44.3 +04:00:22.3 0.02564 −0.01367 8.17 G0 19.22 4.59 0.619 −0.072 1.238 0 HD 126244 14:22:29.7 +55:52:12.4 −0.23388 0.16522 8.22 G0 24.37 5.15 0.647 −0.089 −0.61 0.0 5571 0.751 1 HD 126551 14:23:43.4 +61:56:13.4 0.02858 −0.00647 8.31 F8 12.07 3.72 0.480 −0.004 −0.26 4.3 6297 2.592 1 HD 126511 14:24:49.0 +41:16:30.4 −0.15589 −0.20960 8.36 G5 24.84 5.34 0.757 −0.166 0.14 0.0 5529 0.677 2 HD 126675 14:25:15.7 +53:35:05.8 0.08458 −0.07432 7.11 F8 23.59 3.97 0.526 −0.023 −0.03 2.4 6225 2.095 1 HD 126532 14:25:17.4 +32:47:28.0 0.08826 −0.09342 8.47 G0 31.00 5.93 0.853 −0.244 −0.08 0.0 5155 0.422 1 HD 126583 14:25:58.2 +13:25:59.8 −0.29771 −0.12572 8.06 G5 29.28 5.39 0.750 −0.161 0.01 0.0 5497 0.643 1 HD 126677 14:27:01.2 +01:08:34.5 0.07432 −0.05104 8.02 G0 11.89 3.40 0.686 −0.115 3.855 0 HD 126945 14:28:15.1 +16:07:07.0 0.08146 −0.18083 7.43 G5 15.60 3.40 0.710 −0.132 0.34 2.7 5757 3.916 6 HD 127667 14:32:04.6 +18:50:09.9 −0.24035 0.01344 7.81 F8 16.95 3.96 0.509 −0.016 −0.51 3.5 6089 2.101 2 HD 127852 14:32:24.9 +44:47:05.8 0.05778 −0.15801 8.19 F8 13.25 3.80 0.544 −0.032 −0.12 4.9 6124 2.471 1 HD 128219 14:35:06.0 +23:30:56.3 −0.06925 −0.06650 8.41 G5 14.76 4.26 0.652 −0.092 −0.03 0.9 5792 1.709 1 HD 128660 14:36:55.8 +42:49:52.2 −0.09188 −0.05592 6.59 F8 16.94 2.73 0.505 −0.014 −0.15 6.3 6251 6.510 1 HD 129425 14:40:10.6 +57:42:47.5 −0.09484 0.01167 7.88 F8 19.80 4.36 0.511 −0.017 −0.10 4.2 6250 1.455 1 HD 129171 14:40:18.4 +30:26:37.8 0.09506 −0.04425 7.70 G0 17.98 3.97 0.637 −0.083 0.18 2.6 5928 2.214 7 HD 129209 14:40:28.4 +30:31:13.8 0.09559 −0.04478 7.98 G2 IV 19.82 4.47 0.620 −0.073 0.05 1.7 5932 1.384 1 HD 129357 14:41:22.4 +29:03:31.7 0.01176 −0.18350 7.81 G2 V 21.22 4.44 0.642 −0.086 0.05 1.7 5858 1.440 1 HD 129290 14:41:28.8 +13:36:05.3 −0.38765 −0.07447 8.37 G2 V 14.04 4.11 0.603 −0.063 −0.07 1.5 5940 1.911 1 continued on next page Appendix A. Sample stars N ) ! L (L 1.258 0 1.530 0 3.848 0 0.843 0 0.655 0 1.725 0 0.977 0 1.735 0 1.265 0 3.542 0 2.239 0 eff (K) T i continued on next page sin v (km/s) 0.05 0.4 5603 0.929 1 0.19 3.5 5760 2.847 1 0.10 4.9 6031 4.003 1 0.08 0.7 5632 3.105 1 0.01 6.2 6134 3.757 1 0.26 1.5 5930 1.038 1 0.13 0.0 5461 3.357 1 0.45 0.00.01 5690 3.6 2.148 1 5968 1.370 1 0.34 6.6 6217 3.370 1 0.33 4.4 5106 0.618 1 0.08 0.7 5674 1.475 1 0.32 3.5 6084 4.471 1 0.50 2.2 6007 2.744 1 0.14 4.20.02 6241 1.970 3.2 1 5790 0.594 1 0.09 4.3 6073 1.419 1 0.34 2.8 5735 3.499 1 0.31 0.6 4845 0.287 1 − − − − − − − − − − − − − − − − − − − [Fe/H] 0.130 0.086 0.046 0.069 0.120 0.037 0.051 0.202 0.04 4.9 5349 0.541 1 0.082 0.101 0.040.155 5.3 5779 1.001 1 0.133 0.18 1.5 5685 1.634 5 0.053 0.075 0.030 0.07 4.0 6216 4.786 7 0.045 0.30 4.6 6206 1.288 3 0.151 0.035 0.06 5.0 6177 4.592 3 0.062 0.059 0.18 5.3 6063 1.785 2 0.155 0.080 0.062 0.244 0.03 1.4 5203 0.682 4 0.092 0.260.048 3.7 5915 2.563 3 0.009 0.025 0.02 6.1 6235 2.388 1 0.095 0.228 0.184 0.09 0.0 5438 0.567 2 0.091 0.18 1.7 5885 1.821 17 0.123 0.112 0.103 0.23 3.4 5844 4.104 6 0.026 0.026 0.016 0.094 0.030 0.08 2.6 6220 2.270 1 0.054 0.01 3.0 6021 2.616 1 0.040 0.259 0.07 1.6 5168 0.452 3 0.080 0.035 0.314 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .93 3.35 0.554 .56 4.76 0.583 .07 3.59 0.743 .44 3.83 0.529 .66 5.50 0.834 π 0.56 4.03 0.509 (mas) type Spectral V δ µ (as/yr) 0.16609 8.02 G5 24.38 4.96 0.708 0.12303 8.18 G5 12.70 3.70 0.642 0.11468 8.23 F8 V 10.30 3.29 0.573 0.00673 8.41 K0 17.09 4.57 0.613 0.08457 8.48 K0 10.78 3.64 0.695 0.09651 8.45 K0 15.15 4.35 0.712 0.01734 7.39 G0 15.49 3.34 0.587 0.10954 7.80 G0 Vvar 27.61 5.01 0.624 0.51364 5.04 F8 III-IV 40.46 3.08 0.540 0.12511 7.99 G8 V 29.75 5.36 0.737 0.39230 5.15 F9 IV 39.51 3.13 0.550 0.13112 8.380.06738 G0 V 8.32 14.69 F8 4.22 0.601 14.86 4.18 0.597 0.00735 7.840.03554 8.28 G0 K0 21.16 4.47 26.65 5.41 0.602 0.852 0.06968 7.820.05702 G5 8.26 F9 15.88 V 3.82 15.40 0.651 4.20 0.575 0.24801 8.18 G0 19.12 4.59 0.656 0.28133 8.45 G0 14.09 4.19 0.650 0.02919 8.19 G5 V 11.53 3.50 0.699 0.07705 8.46 F80.02753 7.63 15.72 G0 4.44 0.682 13.72 3.32 0.669 0.15942 7.53 F8 13.30 3.15 0.532 0.08266 8.28 G0 12.49 3.76 0.589 0.05797 7.38 G0 V 25.50 4.41 0.562 0.04222 7.98 K2 V 37.93 5.87 0.870 0.22568 7.87 K0 51.20 6.42 0.931 0.00216 8.44 F8 12.39 3.91 0.550 − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.08255 0.06137 0.05542 0.21742 0.10553 7.30 G5 46.05 5.62 0.803 0.38751 0.00280 8.18 G5 17.31 4.37 0.635 0.22523 0.10778 6.57 G8 IV-V+... 45.32 4.85 0.665 0.13766 0.03742 0.10847 0.05633 8.48 G0 V 16.18 4.52 0.570 0.55421 0.58929 0.28227 7.71 G7 V 27.79 4.93 0.743 0.40535 0.08073 6.64 G0 29.67 4.00 0.631 0.15297 0.014960.51899 0.01166 8.38 F8 10.26 3.44 0.493 0.03476 0.00240 0.16778 0.09451 8.15 G5 30.18 5.55 0.780 0.01450 0.03046 0.20549 0.10953 0.05717 0.17124 8.29 G0 11.96 3.68 0.533 0.00477 0.03424 8.36 G5 V 25.68 5.41 0.654 0.10130 0.03045 8.41 G5 12.48 3.89 0.540 0.02575 0.03885 8.32 G5 10.70 3.47 0.631 0.08704 0.23295 − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 138919 15:34:50.3 +06:48:13.8 HD 139018 15:35:33.9 +06:03:27.1 0.14533 HD 139007 15:34:49.4 +24:59:56.8 0.07497 HD 138820 15:33:41.9 +23:26:38.2 HD 138339HD 138666 15:30:05.5 +39:54:32.5 15:32:57.0 +18:35:49.0 0.02832 0.01565 7.63 F8 13 HD 139213 15:24:28.5 +81:53:42.1 0.19790 0.03448 7.72 G0 25 HD 139813 15:29:24.3 +80:27:00.9 HD 136924 15:22:43.5 +16:15:40.7 HD 138594 15:31:44.3 +39:40:43.9 0.05109 0.05864 7.70 K0 15 HD 139777 15:29:11.9 +80:26:54.9 HD 136902 15:22:35.2 +20:23:48.1 HD 136749 15:19:18.8 +61:22:29.8 BD +332562 15:15:38.3 +33:19:15.3 0.08389 HD 136202 15:19:18.8 +01:45:55.5 0.37324 HD 135633A 15:15:27.2 +22:32:57.4 HD 136064 15:14:38.1 +67:20:48.2HD 136274 0.22151 15:18:59.1 +25:41:30.1 HD 135145 15:12:47.7 +27:55:35.6HD 135906 0.00963 15:16:53.9 +25:45:14.2 0.00875 HIP 74434B 15:12:43.8 +19:17:32.8 HD 133621 15:00:27.5 +71:45:55.6 HD 135143 15:08:30.2 +72:22:09.6 0.08404 HD 134353 15:08:52.0 +19:02:19.1 HD 133091HD 134113 15:02:32.5 +01:01:52.4 HD 135045 15:07:46.5 +08:52:47.2 15:08:35.8 +70:08:48.6 0.00932 0.01059 8.03 G5 14 HD 132505 14:58:58.2 +17:19:37.1 HD 133600 15:05:13.3 +06:17:23.6 HD 132425 14:58:55.3 +06:36:17.5 0.06366 0.10602 8.45 K0 25 HD 132307 14:58:10.9 +06:54:17.8 HD 132406 14:56:54.7 +53:22:55.8 HD 131651 14:54:51.1 +01:06:52.0 HD 130786 14:49:01.2 +36:04:04.6 HD 131526 14:52:20.9 +48:40:14.7 HD 130987 14:49:35.2 +46:34:28.2 0.09569 HD 130396 14:47:31.9HD 131359 +19:03:00.1 14:49:53.6 0.00652 +65:32:50.1 0.00173 7.46 F8 V 2 HD 131179 14:51:53.4 +02:00:53.5 HD 130385 14:47:24.2 +17:45:26.2 HD 130838 14:50:00.6 +12:20:07.8 0.01776 HD 130125 14:46:16.3 +03:30:14.6 0.01696 HD 130215 14:46:03.1 +27:30:44.4 0.01843 HD 130288 14:44:22.0 +63:06:07.4 continued from previous page Star ID HD 129564HD 130004 14:41:47.7 +43:07:48.8 14:45:24.2 +13:50:46.7

154 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 139800 15:39:51.9 +08:31:36.4 0.04448 −0.03376 8.40 F8 16.58 4.50 0.581 −0.051 1.319 0 HD 139839 15:40:21.4 +01:28:07.6 −0.02037 −0.11239 8.11 G0 15.44 4.05 0.700 −0.124 2.136 0 HD 140209 15:42:05.7 +10:58:38.8 −0.17339 −0.02711 8.27 G0 15.47 4.22 0.656 −0.095 1.778 0 HD 140324 15:42:54.5 +04:06:07.2 −0.04012 −0.01262 7.41 G0 18.10 3.70 0.602 −0.062 −0.43 3.9 5796 2.785 1 HD 140514 15:43:10.5 +21:42:05.2 −0.03621 −0.01815 8.50 G2 IV 12.02 3.90 0.684 −0.113 0.03 1.7 5713 2.428 1 HD 140571 15:43:14.3 +31:23:04.4 0.05926 −0.13185 7.26 F8 13.69 2.94 0.455 0.004 −0.19 4.6 6418 5.277 1 HD 140538 15:44:01.8 +02:30:54.6 −0.04496 −0.14473 5.86 G5 V 68.16 5.03 0.684 −0.113 0.09 1.8 5738 0.857 1 HD 140750 15:45:15.2 +04:42:29.6 −0.00992 −0.14144 8.03 G0 14.66 3.86 0.661 −0.099 −0.36 2.0 5627 2.487 1 HD 141024 15:46:22.6 +13:43:14.7 −0.08270 −0.00329 7.93 G5 13.67 3.61 0.685 −0.114 0.03 1.7 5710 3.174 1 HD 141399 15:46:53.8 +46:59:10.6 −0.10824 0.00620 7.21 K0 27.52 4.41 0.770 −0.176 0.29 0.7 5551 1.608 1 HD 142626 15:53:04.5 +53:52:15.7 0.11462 −0.17772 8.28 K0 33.54 5.91 0.835 −0.229 −0.14 0.0 5182 0.424 1 HD 142441 15:54:01.7 +18:46:30.3 −0.14474 −0.08249 8.14 F8 13.76 3.83 0.508 −0.015 2.366 0 HD 142637 15:54:54.9 +22:08:33.1 0.04867 −0.00749 8.22 G0 15.15 4.12 0.614 −0.069 1.904 0 HD 142925 15:54:60.0 +49:51:37.9 −0.05781 0.06727 8.34 K0 13.21 3.94 0.589 −0.054 0.23 3.1 6111 2.216 3 HD 144061 15:58:21.1 +70:53:40.0 −0.05970 0.25389 7.26 G5 34.35 4.94 0.654 −0.094 −0.28 0.1 5683 0.915 1 HD 143292 15:59:01.9 +18:19:33.1 −0.05981 0.00286 7.79 F8 15.13 3.69 0.517 −0.019 −0.33 4.4 6134 2.701 1 HD 143436 16:00:18.8 +00:08:13.2 −0.13583 −0.11031 8.05 G0 23.04 4.86 0.643 −0.087 0.02 2.1 5843 0.979 1 HD 144031 16:01:03.1 +55:27:12.2 −0.00154 −0.01379 8.43 F8 14.53 4.24 0.497 −0.011 −0.05 3.7 6321 1.616 1 HD 143893 16:01:45.7 +31:52:34.4 0.01052 0.01210 7.67 G0 13.10 3.26 0.597 −0.059 −0.46 3.7 5800 4.165 1 HD 144302 16:01:51.8 +59:37:50.6 −0.12432 0.05382 7.85 G5 13.81 3.55 0.667 −0.102 0.23 4.4 5850 3.317 3 HD 144270 16:03:42.8 +34:00:46.7 0.05521 0.02976 8.21 F8 21.39 4.86 0.658 −0.097 −0.16 3.2 5719 0.988 1 155 HD 144172 16:04:29.6 +00:40:20.0 0.00489 −0.08266 6.81 F8 19.25 3.23 0.451 0.006 −0.62 4.1 6257 4.033 1 HD 144705 16:05:52.5 +35:00:21.7 −0.06066 −0.25228 8.04 G0 17.10 4.20 0.569 −0.044 −0.60 0.9 5839 1.728 1 HD 145046 16:06:19.3 +53:30:54.2 0.07020 −0.09380 8.35 G5 13.57 4.01 0.580 −0.050 0.00 2.3 6048 2.070 1 HD 145224 16:08:27.0 +33:43:59.1 −0.05439 0.01479 8.37 G5 18.31 4.68 0.644 −0.087 0.09 2.0 5868 1.156 1 HD 145514 16:09:41.2 +40:07:08.6 −0.15002 0.26349 8.31 G0 18.34 4.63 0.676 −0.108 −0.13 0.0 5673 1.234 1 HD 146044 16:10:47.3 +59:40:31.5 −0.03843 −0.11520 8.12 G0 24.45 5.06 0.688 −0.116 −0.12 0.0 5638 0.836 1 HD 145645 16:10:54.1 +30:59:07.5 −0.08428 −0.17172 7.42 G5 18.44 3.75 0.706 −0.129 0.12 1.6 5679 2.829 2 BD +472312 16:11:08.4 +47:32:32.4 0.02470 0.03650 8.31 G5 25.70 5.36 0.714 −0.135 0.646 0 HD 146079 16:12:15.2 +46:13:33.7 −0.12239 0.22050 7.80 F8 13.70 3.48 0.558 −0.038 3.336 0 HD 146099 16:12:43.2 +43:38:58.5 0.11844 0.00276 8.12 F8 12.27 3.56 0.524 −0.023 3.056 0 HD 146735 16:14:44.6 +57:01:34.6 0.06570 −0.05020 8.39 G0 12.70 3.91 0.639 −0.084 0.05 2.0 5868 2.342 1 HD 146868 16:14:57.0 +60:40:11.1 0.02507 0.43834 7.67 G5 32.72 5.24 0.659 −0.097 0.696 0 HD 146588 16:16:14.6 +19:31:29.2 0.03171 0.26571 7.80 G0 22.10 4.52 0.522 −0.022 −0.68 2.5 5972 1.261 1 HD 147187 16:16:53.1 +60:32:05.6 0.10760 0.00310 8.20 G5 16.65 4.31 0.766 −0.173 0.18 1.4 5518 1.759 22 HD 146644 16:17:01.1 +07:09:06.3 −0.12234 −0.05789 8.32 G5 16.44 4.40 0.697 −0.122 0.07 1.0 5688 1.545 1 HD 147062 16:19:25.9 +05:32:17.2 −0.03247 −0.02443 7.59 G0 15.30 3.51 0.626 −0.077 3.364 0 HD 147528 16:22:07.0 +12:12:52.4 0.06623 −0.06884 8.18 G0 19.73 4.66 0.544 −0.032 −0.19 4.3 6095 1.119 1 HD 147750 16:23:06.0 +17:28:07.7 −0.13162 0.30421 8.44 G0 24.97 5.43 0.724 −0.142 −0.09 0.0 5536 0.609 1 BD +342776 16:23:14.7 +33:41:49.5 −0.03913 0.10185 7.65 F8 16.27 3.71 0.495 −0.010 −0.20 5.6 6267 2.630 1 HD 147887 16:24:11.3 +07:27:33.0 −0.00032 −0.09957 7.99 F8 12.53 3.48 0.637 −0.083 0.09 2.8 5891 3.477 3 HD 148164 16:25:44.5 +11:55:08.6 −0.08993 0.00167 8.23 F8 13.39 3.86 0.589 −0.054 0.13 5.4 6070 2.386 4 HD 148492A 16:27:46.5 +20:53:43.2 0.03296 −0.11749 8.24 G 10.92 3.43 0.600 −0.061 3.568 0 HD 149222 16:29:01.2 +64:47:12.2 −0.14063 0.11337 7.68 G0 16.22 3.73 0.585 −0.052 0.18 2.3 6104 2.684 5 HD 148816 16:30:28.4 +04:10:41.6 −0.43273 −1.39234 7.27 F8 V 24.34 4.20 0.545 −0.033 −0.90 2.0 1.711 1 HD 148998 16:31:34.4 +08:17:58.3 0.04035 0.07201 8.26 G5 12.86 3.81 0.602 −0.062 −0.00 2.6 5972 2.517 1 continued on next page Appendix A. Sample stars N ) ! L (L 1.640 0 0.891 0 1.701 0 2.390 0 1.384 0 1.204 0 1.440 0 eff (K) T i continued on next page sin v (km/s) 0.60 2.0 5832 2.061 1 0.28 1.0 4630 0.266 1 0.25 0.0 5624 0.877 1 0.14 0.9 4697 0.283 1 99.9 99.9 2.979 1 0.20 0.0 5301 2.375 1 0.15 1.6 4884 0.297 1 0.04 1.0 5729 1.252 1 0.20 4.2 6177 1.557 1 0.03 3.30.44 6304 5.0 2.059 1 6279 3.339 1 0.06 0.0 5625 0.768 1 0.010.62 2.50.78 4.2 5927 1.6 3.798 6109 2.554 1 1 1.893 1 0.01 3.5 6110 4.920 1 0.19 2.7 4704 0.235 5 0.16 2.2 5693 0.668 1 0.50 1.8 5849 1.116 1 0.08 0.0 5792 1.392 1 0.81 0.5 1.108 1 0.30 1.8 6011 1.848 1 − − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.112 0.10 5.4 5751 1.078 1 0.077 0.124 0.04 0.0 5666 2.739 1 0.115 0.045 0.057 0.186 0.402 0.108 0.398 0.035 0.189 0.156 0.32 0.7 5643 1.522 4 0.323 0.073 0.104 0.021 0.201 0.08 0.9 5371 0.705 3 0.014 0.258 0.14 0.4 5200 0.672 8 0.124 0.069 0.008 0.033 0.089 0.04 4.1 5841 4.321 1 0.176 0.300.040 1.3 5555 3.037 4 0.386 0.072 0.102 0.108 0.01 2.8 5730 3.760 1 0.087 0.10 3.6 5872 2.824 1 0.049 0.068 0.12 4.6 5988 3.591 2 0.089 0.051 0.03 3.4 6050 2.808 1 0.117 0.12 1.8 5731 1.581 3 0.056 0.0300.068 0.03 0.29 4.5 4.5 6196 1.803 6061 2.287 1 26 0.146 0.09 1.0 5592 3.788 2 0.170 0.21 1.1 5539 0.794 2 0.051 0.1110.037 0.20 3.8 5796 2.468 14 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .17 3.42 0.676 .72 5.29 0.666 π 6.79 5.16 0.700 (mas) type Spectral V δ µ (as/yr) 0.00580 8.31 G5 19.71 4.78 0.681 0.10353 7.57 G5 17.45 3.78 0.700 0.30865 7.94 G0 V 18.13 4.23 0.594 0.19516 7.93 K2 53.97 6.59 1.019 0.04763 8.50 F8 Vvar 10.47 3.60 0.549 0.18448 7.22 G8 V 22.68 4.00 0.786 0.02994 8.41 G0 V 16.27 4.47 0.620 0.11755 7.95 K0 48.69 6.39 0.941 0.15118 8.43 K0 17.22 4.61 0.670 0.02539 7.99 F8 15.77 3.98 0.504 0.04941 8.300.21677 G5 7.28 10.32 G0 3.37 23.04 0.614 4.09 0.545 0.17843 6.52 G5 IV:0.13487 7.00 27.58 F8 IV 3.72 0.770 16.29 3.06 0.561 0.01657 8.38 F8 10.21 3.43 0.612 0.28310 8.06 G0 19.24 4.48 0.646 0.24324 7.95 K00.06671 19.26 7.72 4.37 G0 V 0.690 15.57 3.68 0.583 0.10777 6.480.08442 8.39 F8 V G0 34.00 4.14 12.76 0.541 3.92 0.611 0.19210 7.50 G5 V 15.46 3.45 0.730 0.22600 8.05 G5 26.56 5.17 0.763 0.07803 8.32 F8 18.75 4.69 0.582 0.08503 7.890.46527 7.10 G0 F8 V 15.75 25.37 3.88 4.12 0.680 0.555 − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.36174 0.07409 8.48 G1 V 14.51 4.29 0.626 0.07897 0.03290 0.08505 8.03 G5 24.64 4.99 0.687 0.16501 0.26980 6.95 F9 V 25.88 4.01 0.571 0.02453 0.25534 0.09330 8.14 G5 14.79 3.99 0.782 0.00197 0.11174 8.05 G5 24.55 5.00 0.676 0.07225 0.21639 8.43 K2 41.58 6.52 1.015 0.15772 0.11806 8.04 K0 19.13 4.45 0.744 0.08598 0.14879 0.06909 7.48 F8 22.97 4.29 0.520 0.00200 0.16096 8.07 K0 28.27 5.33 0.801 0.20260 0.23548 8.33 K2 V 26.38 5.44 0.869 0.02107 0.04383 0.06296 8.41 F8 11.64 3.74 0.491 0.02471 0.14102 6.890.03160 F8 V0.22350 18.72 3.25 0.646 0.16644 0.072320.06521 0.02223 8.38 G0 10.26 3.44 0.465 0.001 0.13909 0.26249 7.93 K0 57.13 6.71 1.004 0.19001 0.08603 8.14 G5 19.75 4.62 0.618 0.05161 0.11047 8.03 G0 13.68 3.71 0.644 0.13033 0.08593 8.22 G5 19.60 4.68 0.578 0.05113 0.09514 0.00218 7.67 F8 22.31 4.41 0.591 0.00056 0.03301 0.20953 0.01006 0.00252 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 158415 17:27:13.9 +42:13:04.9 0.07577 HD 158226 17:26:43.2 +31:04:37.8 HD 158096 17:24:37.2 +50:41:08.7 HD 157637 17:23:53.0 +13:05:58.2 HD 157089 17:21:07.1 +01:26:35.0 HD 156968 17:20:11.6 +09:27:39.4 HD 156893 17:19:46.4 +06:33:13.4 HD 156985 17:17:50.4 +52:26:49.6 0.01727 HD 156728 17:16:39.2 +50:36:23.3 HD 156668 17:17:40.5 +29:13:38.0 HD 155937 17:13:57.8 +16:21:00.6HD 157102 0.02624 17:16:59.0 +62:34:41.3 HD 156111 17:14:57.1 +19:40:57.4 BD +292963 17:14:11.6 +29:34:09.6 0.02142 HD 155712 17:12:37.7 +18:21:04.3 0.10178 HD 156558 17:12:29.9 +69:18:15.9 0.01209 HD 153970 17:00:28.9 +45:46:43.7 HD 156279 17:12:23.2 +63:21:07.5 HD 155456 17:10:37.6 +24:31:56.8 HD 153701HD 153915 17:00:28.6 +15:09:33.7 17:01:55.4 +07:39:30.8 0.05177 0.12956 8.02 G5 V 2 HD 154714 17:06:21.9 +20:13:39.7 HD 153376 16:58:37.9 +15:27:15.7 HD 154064 17:02:42.7 +12:56:32.2 HD 155358 17:09:34.7 +33:21:21.1 HD 154160HD 155952 17:03:10.4 +14:30:40.8 HD 155193 17:09:02.2 +69:06:15.2 17:09:41.1 +10:02:19.3 HD 153525 16:57:42.2 +47:21:43.7 HD 154159 16:57:13.6 +71:27:47.7 HD 153205 16:54:03.1 +61:38:47.7 0.02344 0.03423 8.18 G0 11 HD 152012 16:47:41.5 +56:53:33.0 0.01762 0.04471 8.41 G0 23 HD 152305 16:51:15.3 +34:17:21.8 HD 151252 16:41:53.6 +63:48:56.3 HD 151426HD 151689 16:44:21.1 +54:54:59.9 16:46:30.7 +47:26:57.0 0.14411 HD 150633 16:41:50.2 +13:08:55.9 0.00430 HD 152264 16:51:12.9 +29:34:15.9 0.09477 HD 151044 16:42:27.8 +49:56:11.2 0.13343 HD 150554 16:40:56.4 +21:56:53.3 HD 151501 16:46:14.4 +36:41:20.0 HD 150205 16:38:16.2 +29:40:20.9 HD 149933 16:37:45.9 +00:02:24.6 HD 149380 16:33:21.1 +25:41:22.9 0.23555 continued from previous page Star ID HD 149143HD 149890 16:32:51.1 +02:05:05.4 16:36:26.0 +30:56:30.0

156 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 158332 17:27:34.5 +26:47:41.8 −0.10127 0.27533 7.69 K1 IV 33.36 5.31 0.820 −0.216 0.20 0.2 5365 0.728 14 HD 158573 17:29:12.8 +22:07:56.7 −0.03418 −0.12149 8.00 G5 11.27 3.26 0.669 −0.103 −0.05 3.9 5729 4.337 1 HD 159265 17:29:16.6 +63:52:08.9 0.00031 −0.18138 8.42 G0 22.18 5.15 0.708 −0.130 −0.06 1.1 5599 0.780 1 HD 159329 17:29:44.4 +63:51:09.6 0.00248 −0.18216 7.66 F9 V 22.32 4.40 0.581 −0.051 −0.05 3.3 6024 1.447 1 HD 158736 17:30:50.3 +01:34:41.7 −0.02682 0.11126 7.93 F8 16.28 3.99 0.527 −0.024 2.059 0 HD 159333 17:33:52.8 +08:06:13.6 −0.01729 −0.05761 7.85 G0 10.59 2.97 0.596 −0.059 5.440 0 HD 159356 17:33:56.1 +08:09:57.8 −0.01287 −0.05916 8.50 G 10.09 3.52 0.656 −0.095 3.388 0 HD 159482 17:34:43.0 +06:00:51.6 −0.47864 0.37373 8.37 G0 V 20.90 4.97 0.577 −0.049 0.854 0 HD 160933 17:36:40.0 +69:34:14.6 −0.05374 −0.21754 6.33 F9 V 23.53 3.19 0.596 −0.059 −0.28 4.0 5878 4.442 1 HD 160539 17:36:42.4 +59:30:08.1 −0.04385 −0.08978 8.29 G0 11.49 3.59 0.565 −0.042 −0.02 4.2 6092 3.026 1 HD 160013 17:36:55.7 +19:55:20.5 −0.10325 −0.08322 8.50 G5 24.05 5.41 0.764 −0.171 0.19 0.8 5528 0.637 4 HD 160216 17:37:19.2 +31:11:26.7 0.00801 0.03112 8.37 G5 13.50 4.02 0.594 −0.057 −0.18 5.9 5926 2.064 1 HD 160135 17:37:26.8 +22:21:11.8 −0.01148 0.04879 8.38 G0 17.27 4.57 0.623 −0.075 −0.07 4.1 5872 1.265 1 HD 160420 17:37:55.2 +41:00:39.1 0.12164 0.00919 8.05 G5 18.31 4.36 0.583 −0.051 −0.20 2.5 5955 1.501 1 HD 160508 17:39:12.7 +26:45:27.1 0.03497 −0.06575 8.11 F8 V 10.53 3.22 0.543 −0.032 −0.02 4.0 6169 4.215 1 HD 161237 17:39:41.4 +64:32:11.1 −0.10363 −0.08511 8.37 G5 12.22 3.81 0.697 −0.122 0.05 1.9 5679 2.660 1 HD 161284 17:39:55.7 +65:00:05.9 −0.02126 0.10642 8.39 K0 37.83 6.28 0.930 −0.313 −0.01 2.6 4970 0.326 1 HD 161897 17:41:06.8 +72:25:13.2 −0.12186 0.29722 7.60 K0 34.57 5.29 0.720 −0.139 0.01 1.0 5590 0.691 1 HD 161194 17:41:56.1 +45:29:45.5 −0.07279 −0.15857 8.45 G0 12.34 3.91 0.528 −0.024 2.216 0 BD +442753 17:41:58.3 +44:44:21.3 −0.12697 0.02848 8.47 G0 13.35 4.10 0.725 −0.142 2.074 0 HD 162195 17:43:55.8 +69:04:26.4 −0.06179 −0.04321 7.82 F8 15.80 3.81 0.554 −0.037 0.01 3.9 6142 2.459 1 157 HD 161479 17:45:02.9 +19:17:25.6 0.15682 0.05010 8.14 K0 21.63 4.82 0.769 −0.175 0.19 3.9 5513 1.102 4 HD 161728 17:46:27.8 +17:34:08.4 −0.03428 −0.17150 8.03 G0 15.88 4.03 0.631 −0.080 −0.06 1.7 5850 2.089 2 HD 162209 17:48:13.0 +38:13:57.3 0.05889 −0.07703 7.78 G0 18.90 4.16 0.647 −0.089 0.02 1.4 5830 1.869 1 HD 162197 17:48:20.4 +33:48:01.0 −0.06990 −0.04259 8.18 G0 17.91 4.45 0.632 −0.080 −0.31 1.3 5744 1.419 1 HD 162735 17:51:53.0 +13:33:07.0 −0.02518 −0.05816 8.27 G5 10.92 3.46 0.720 −0.139 −0.15 2.6 5524 3.729 1 BD +213245 17:53:30.0 +21:19:31.0 −0.07150 0.05734 8.50 K0 40.22 6.52 0.940 −0.322 0.264 0 HD 163714 17:53:36.5 +61:02:37.9 0.00791 0.02674 7.88 F8 19.42 4.32 0.599 −0.060 −0.18 3.0 5908 1.570 1 HD 163607 17:53:40.6 +56:23:31.1 −0.07469 0.12056 8.00 G5 14.40 3.79 0.777 −0.182 0.11 0.0 5456 2.863 2 HD 164057 17:54:38.6 +64:59:34.5 0.00456 −0.05916 8.21 F8 12.55 3.70 0.508 −0.015 −0.07 3.6 6273 2.667 1 HD 163589 17:54:40.0 +45:33:14.9 0.04127 0.02498 7.68 G5 15.59 3.64 0.727 −0.144 0.19 1.3 5642 3.174 18 HD 163620 17:55:07.9 +41:27:56.1 −0.02917 −0.17686 8.21 G0 14.30 3.99 0.555 −0.037 −0.33 2.3 5999 2.084 1 HD 163288 17:55:21.0 +00:40:01.2 −0.01110 −0.01032 8.30 G0 14.01 4.03 0.529 −0.025 −0.20 13.0 6144 1.986 1 HD 163441 17:55:45.3 +11:48:14.1 0.00209 −0.03319 8.43 G5 17.80 4.68 0.685 −0.114 −0.03 1.7 5685 1.185 1 HD 163622 17:56:42.0 +11:52:19.4 0.03417 −0.13937 8.12 G5 18.71 4.48 0.606 −0.065 −0.02 3.3 5950 1.361 1 HD 164509 18:01:31.2 +00:06:16.4 −0.00578 −0.02039 8.10 G5 19.30 4.53 0.665 −0.101 0.11 2.8 5807 1.344 2 HD 164651 18:02:17.9 +00:06:15.1 0.17817 0.05280 7.66 G5 30.61 5.09 0.746 −0.158 −0.12 0.0 5456 0.846 1 HD 165173 18:04:26.4 +06:26:11.0 −0.06357 0.07102 7.96 G5 30.59 5.39 0.762 −0.170 0.649 0 HD 165504 18:04:52.4 +33:16:20.4 0.00580 −0.19753 7.67 G0 21.55 4.34 0.639 −0.084 0.01 2.7 5852 1.576 1 HD 165476 18:05:52.3 +07:05:26.3 −0.20928 −0.10339 7.65 G5 22.16 4.38 0.624 −0.075 1.507 0 HD 165626 18:06:20.2 +15:34:31.3 −0.01924 −0.20910 8.36 K0 12.53 3.85 0.640 −0.085 −0.29 1.4 5725 2.477 1 HD 165672 18:06:48.8 +06:24:38.0 −0.00449 −0.04363 7.77 G5 22.82 4.56 0.663 −0.100 1.306 0 HD 166579 18:07:36.6 +58:58:32.4 0.00148 0.13039 8.43 G5 13.86 4.14 0.560 −0.039 0.07 4.8 6146 1.818 1 HD 166685 18:11:23.8 +11:31:43.4 −0.02470 −0.01624 8.40 G5 15.48 4.35 0.520 −0.021 1.474 0 HD 167081 18:12:43.3 +24:23:36.4 −0.12681 −0.03313 8.20 F8 20.32 4.74 0.697 −0.122 1.129 0 HD 167215 18:12:59.4 +28:15:27.4 −0.00286 0.11374 8.08 F8 12.72 3.60 0.578 −0.049 3.017 0 continued on next page Appendix A. Sample stars N ) ! L (L 0.618 0 0.632 0 1.911 0 2.289 0 1.986 0 0.959 0 3.113 0 2.144 0 2.747 0 eff (K) T i continued on next page sin v (km/s) 0.46 5.2 6038 2.692 1 0.41 3.9 5883 3.631 1 0.10 3.2 6139 2.0910.42 1 1.2 5742 2.017 1 0.27 4.2 62160.12 3.139 0.7 1 0.20 5490 2.6 1.852 6038 1 1.583 1 0.51 6.1 6053 3.407 1 0.05 1.6 5884 1.264 1 0.13 1.7 5808 0.957 1 0.28 1.3 5162 0.369 1 0.13 2.4 4568 0.235 8 0.25 4.8 5969 3.858 1 0.10 2.7 6052 2.542 1 0.56 7.0 6079 3.644 1 0.04 3.6 5762 1.328 1 0.61 2.5 5846 1.532 1 0.09 3.3 5922 1.072 1 0.31 2.6 6021 2.716 1 0.13 0.6 5504 2.061 1 − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.025 0.050 0.075 0.11 3.5 5946 2.818 4 0.031 0.072 0.177 0.171 0.039 0.083 0.250.012 4.7 59610.149 3.709 3 0.021 0.053 0.074 0.059 0.085 0.082 0.051 0.25 4.0 6147 3.736 3 0.055 0.072 0.11 2.5 5963 0.984 3 0.051 0.25 6.0 6147 4.703 7 0.089 0.06 4.5 5849 1.317 2 0.218 0.459 0.043 0.046 0.077 0.050.139 2.10.058 0.10 0.20 5912 2.0 1.199 6.1 5627 1 3.401 6078 3.186 1 3 0.028 0.096 0.13 4.0 5842 3.788 2 0.043 0.077 0.103 0.10 2.1 5790 1.680 1 0.014 0.098 0.080 0.04 3.8 5891 3.221 1 0.043 0.161 0.220.065 1.6 5583 1.307 16 0.203 0.15 2.3 5391 0.588 6 0.293 0.12 1.8 5083 0.399 6 0.106 0.15 3.6 5798 4.682 7 0.035 0.145 0.111 0.07 1.4 5742 1.038 1 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .21 5.45 0.771 .68 4.29 0.559 .71 3.44 0.519 .58 4.57 0.622 .85 4.54 0.646 .14 6.05 0.822 .48 3.56 0.720 .39 4.63 0.626 .86 3.95 0.536 .61 4.33 0.566 .60 4.74 0.606 .50 3.70 0.551 π 0.00 5.53 0.804 0.02 4.82 0.680 21.40 4.62 0.750 (mas) type Spectral V δ µ (as/yr) 0.24017 6.510.16118 F8 V 7.85 23.84 G5 3.40 0.579 14.80 3.70 0.623 0.03759 7.65 F8 18.49 3.98 0.542 0.46884 8.32 G2 V 26.30 5.42 0.764 0.05155 7.61 F80.03087 8.04 15.23 3.52 G5 0.501 17.32 4.23 0.735 0.06347 7.79 G00.03236 8.44 16.76 3.91 G5 0.597 0.12208 13.46 7.51 4.09 F8 0.640 14.86 3.37 0.581 0.05747 7.58 G3 Vvar 28.27 4.84 0.618 0.09340 7.64 F8 12.48 3.12 0.581 0.43685 7.97 K0 57.79 6.78 1.070 0.06387 7.89 G0 13.60 3.56 0.566 0.11034 8.00 F8 11.620.03626 3.33 8.32 0.573 G0 11.13 3.55 0.595 0.10424 7.84 G0 15.100.08391 7.90 3.73 0.626 G0 18.94 4.29 0.669 0.09173 8.20 G0 10.78 3.36 0.506 0.008720.14473 8.38 8.46 G5 G0 17.06 10.46 4.54 3.56 0.660 0.631 0.17311 8.06 K0 39.43 6.04 0.908 0.08873 7.91 G0 IV 11.33 3.18 0.673 0.04828 7.07 G5 25.64 4.11 0.729 − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.00312 0.11066 8.23 F8 12.39 3.70 0.529 0.04290 0.34879 0.19992 8.27 G1 V 14.37 4.06 0.619 0.01359 0.00181 7.33 G0 16.48 3.41 0.636 0.19597 0.13740 0.14871 8.34 G5 14.21 4.10 0.587 0.09519 0.00214 0.07737 8.26 K0 21.10 4.88 0.635 0.06246 0.04890 8.43 G0 19.27 4.85 0.590 0.05154 0.06939 0.01174 0.00335 8.26 G5 10.66 3.40 0.657 0.01250 0.06265 8.37 G5 12.09 3.78 0.567 0.05256 0.04576 0.01565 0.00494 − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) continued from previous page Star ID HD 167216 18:13:02.3 +28:14:47.9 HD 167588 18:14:44.0 +29:12:26.3 0.01179 HD 168960 18:21:50.4 +15:37:33.7 HD 168093 18:15:20.0 +50:58:16.6 0.01554 HD 170357 18:26:37.8 +46:05:01.5 HD 168603 18:19:08.8 +33:13:52.6 0.00487 0.13731 8.28 K0 27 HD 172083 18:31:03.0 +73:05:01.7 0.06231 0.10264 8.18 G0 16 HD 169925 18:15:34.5HD 169244 +79:03:53.6 18:22:08.9HD 172339 +36:18:06.3 0.04858 18:27:33.2 +79:47:41.6 0.07788 0.04536 8.29 G5 10 HD 169889 18:26:22.0 +08:36:56.8 HD 171009 18:30:57.5 +33:55:58.8 HD 170294 18:28:34.4 +00:47:49.4 HD 172557 18:36:33.9 +63:21:44.8 0.10236 0.16145 8.22 G5 18 HD 169359 18:23:47.1 +14:54:27.8BD +163516 0.12469 18:28:44.5HD 171607 +16:46:26.1 18:31:47.9 0.17392 +61:46:44.7 HD 170096 18:27:37.0 +02:22:59.8 HD 172669 18:36:23.3 +66:54:47.7 0.06973 HD 172085 18:37:10.5 +24:25:59.9 HD 171489 18:34:35.6 +12:32:20.2 0.02585 0.02715 8.28 G5 17 HD 172393 18:38:01.9 +42:39:54.8 0.29437 0.06311 8.32 G5 35 HD 173700 18:43:18.1HD 174080 +57:49:06.5 18:48:29.2 +10:44:43.6 0.12876 HD 171951A 18:37:15.3 +07:31:42.7 0.05643 HD 176373 18:47:32.8 +80:09:37.9 0.02159 0.07034 8.46 G5 10 HD 172718 18:41:21.7HD 173071 +07:54:13.1 0.03717 18:43:15.7 +09:02:28.6HD 175441 0.03881 18:51:38.1 0.03087 +60:26:02.8 8.19 0.00603 G0 19 HD 173091 18:43:19.5 +09:09:46.9 0.08835 0.12198 8.24 F8 13 HD 173289 18:42:27.7 +44:16:24.9 HD 173024 18:41:27.4 +40:28:35.3 HD 172867 18:42:05.0 +11:45:39.6 HD 175425 18:53:38.1 +37:59:06.9 HD 176213 18:58:18.1 +23:50:43.3 0.04965 HD 177780 19:04:16.4 +41:00:11.4 0.01625 HD 177572 19:04:26.3 +22:39:55.8 HD 180841 19:09:38.0 +76:10:58.3 0.15981 0.05329 8.23 F8 16 BD +343438B 19:09:03.1HD 180712 +34:35:59.5 0.06477 19:14:01.4 +59:33:06.6 0.19241 7.97 0.04624 0.11228 G5 7.97 F8 22 HD 180161 19:12:11.2 +57:40:19.2 0.21775 0.40826 7.04 G8 V 5 HD 180263 19:15:35.0 +11:33:17.0 0.18398 HD 180502 19:15:38.1 +29:07:19.0 HD 180683 19:15:48.8 +38:22:49.1 0.02925 0.24224 7.75 G0 15 HD 182189 19:17:22.3 +71:20:49.3 0.00703 HD 181047 19:17:53.7 +25:22:11.0 0.15975 0.23278 8.31 G8 V 2

158 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 182335 19:23:24.7 +20:34:07.6 0.01313 −0.18693 7.91 G2 IVws 21.26 4.55 0.578 −0.049 1.258 0 HD 182407 19:24:23.0 +05:33:52.0 0.00768 0.02773 7.77 G0 V 10.47 2.87 0.617 −0.071 0.07 6.4 5950 6.031 1 HD 182619 19:24:41.9 +22:12:05.5 0.13498 −0.08409 7.81 G5 29.97 5.19 0.718 −0.138 −0.09 0.0 5555 0.757 1 HD 182758 19:24:56.1 +31:48:04.3 0.07832 −0.09497 8.18 G0 15.13 4.08 0.573 −0.046 −0.75 1.9 1.934 1 HD 183298 19:26:22.1 +50:59:14.3 0.02974 0.07371 8.15 G0 11.85 3.52 0.593 −0.057 0.05 3.1 6024 3.272 1 HD 183162 19:27:51.8 +08:58:10.8 0.09629 0.07693 8.06 G0 14.85 3.92 0.606 −0.065 0.07 4.4 5987 2.280 1 HD 183263 19:28:24.6 +08:21:29.0 −0.01778 −0.03314 7.86 G2 IV 18.93 4.25 0.678 −0.110 0.17 2.8 5790 1.754 7 HD 184601 19:31:44.6 +60:51:51.1 −0.05309 −0.09167 8.26 G0 12.76 3.79 0.540 −0.030 −1.15 2.6 2.489 1 HD 183993 19:32:00.9 +05:23:54.9 0.08144 −0.11819 8.44 G8 IV 20.50 5.00 0.805 −0.204 0.25 0.9 5429 0.959 2 HD 184448 19:32:02.5 +50:10:53.6 −0.09099 0.29217 8.05 F8 19.16 4.46 0.659 −0.097 1.428 0 HD 184403 19:33:26.2 +23:29:51.0 −0.03206 −0.00023 7.81 G0 22.52 4.57 0.654 −0.094 0.04 1.1 5815 1.287 1 HD 185239 19:34:51.9 +57:46:21.5 −0.03507 0.02302 8.15 F8 12.21 3.58 0.452 0.005 −0.14 5.7 6450 2.924 6 HD 184592 19:34:55.7 +11:25:26.8 0.27920 0.00756 7.91 G5 23.69 4.78 0.670 −0.104 0.02 1.9 5754 1.071 1 HD 184768 19:36:00.7 +00:05:28.2 −0.01466 −0.38496 7.55 G5 25.56 4.59 0.675 −0.107 −0.01 0.0 5726 1.279 1 HD 184979 19:36:05.5 +24:43:49.1 0.08414 0.04188 8.25 G5 17.28 4.44 0.574 −0.047 −0.20 3.4 5986 1.389 1 HD 185269 19:37:11.8 +28:29:59.5 −0.03213 −0.08081 6.67 G0 IV 21.11 3.29 0.606 −0.065 0.15 5.5 6020 4.074 44 HD 189805 19:41:05.0 +84:43:47.6 0.06733 0.06644 8.48 G0 19.64 4.95 0.647 −0.089 −0.28 1.4 5706 0.903 1 HD 331093 19:42:28.8 +31:44:25.5 −0.01610 −0.19937 8.38 K0 19.60 4.84 0.770 −0.176 0.15 0.0 5493 1.082 1 HD 186413 19:44:04.4 +03:30:27.8 0.05282 −0.12833 7.99 G0 13.48 3.64 0.625 −0.076 −0.08 2.9 5862 2.981 1 HD 187253 19:46:19.2 +54:38:18.4 −0.00757 −0.02771 8.27 F8 15.31 4.19 0.498 −0.011 −0.26 6.1 6231 1.692 1 HD 186843 19:46:36.1 +07:03:20.4 0.04054 −0.00806 8.20 G5 15.74 4.19 0.590 −0.055 0.06 3.0 6038 1.762 6 159 HD 186932 19:46:37.8 +17:48:10.5 0.02182 −0.04572 8.10 G0 11.03 3.31 0.608 −0.066 0.04 3.8 5968 4.003 1 HD 187057 19:47:39.4 +07:38:57.2 −0.0196 −0.04296 8.37 G0 11.35 3.64 0.560 −0.039 0.06 3.2 6142 2.881 1 HD 187792 19:47:46.5 +63:51:37.4 −0.05024 0.03922 7.86 G0 18.80 4.23 0.602 −0.062 −0.07 2.8 5943 1.709 1 HD 187748 19:48:15.4 +59:25:22.4 0.01575 0.11651 6.64 G0 35.25 4.38 0.592 −0.056 0.04 6.4 6023 1.480 2 HD 187876 19:49:12.2 +57:24:33.6 −0.00702 −0.06054 7.76 G0 20.50 4.32 0.613 −0.069 −0.04 2.2 5918 1.583 1 HD 187548 19:49:26.8 +28:36:36.6 0.07546 0.07193 7.98 G0 V 22.20 4.71 0.518 −0.020 −0.55 2.8 6040 1.057 1 HD 187691 19:51:01.6 +10:24:56.6 0.24009 −0.13451 5.12 F8 V 51.57 3.68 0.563 −0.041 0.18 4.0 6181 2.782 3 HD 187881 19:51:32.0 +22:16:50.0 −0.11079 −0.12643 7.90 G5 12.28 3.35 0.685 −0.114 −0.17 2.7 5627 4.033 1 HD 187882 19:51:37.2 +22:09:21.9 0.06066 −0.05573 8.18 G5 13.15 3.77 0.602 −0.062 −0.51 2.3 5763 2.611 3 HD 187944 19:51:59.4 +17:50:18.4 0.04487 −0.05508 8.17 G0 10.92 3.36 0.500 −0.012 −0.34 4.2 6191 3.637 1 HD 188307 19:52:52.6 +41:04:49.8 −0.01697 0.05962 7.72 F8 18.81 4.09 0.572 −0.046 −0.11 4.6 6030 1.916 1 HD 188326 19:53:01.6 +38:46:24.2 −0.04328 0.34204 7.57 G8 IVvar 17.72 3.81 0.773 −0.178 2.800 0 HD 188345 19:53:59.7 +20:00:34.6 0.07400 −0.01310 7.97 G5 14.01 3.70 0.628 −0.078 2.826 0 HD 354088 19:54:47.7 +17:07:08.7 0.00658 −0.10492 8.46 G0 12.94 4.02 0.573 −0.046 −0.08 3.0 6039 2.044 1 HD 189749 19:57:53.9 +62:57:50.5 0.12525 0.13519 8.37 G 17.34 4.57 0.593 −0.057 −0.35 3.0 5859 1.244 1 HD 189733 20:00:43.7 +22:42:39.1 −0.00249 −0.25081 7.67 G5 51.94 6.25 0.932 −0.315 −0.07 3.5 4940 0.336 26 HD 190508 20:02:41.9 +53:53:21.8 −0.09364 −0.11786 8.32 G5 16.25 4.37 0.680 −0.111 0.06 1.2 5738 1.572 1 HD 190470 20:04:10.1 +25:47:24.9 −0.07631 −0.03904 7.82 K3 V 46.28 6.15 0.924 −0.308 0.07 1.5 5019 0.366 3 HD 190605 20:04:55.1 +26:03:15.2 0.05135 −0.38121 7.69 G2 V 22.49 4.45 0.669 −0.103 0.03 0.8 5761 1.449 1 HD 190594 20:05:33.7 +03:30:10.0 0.13743 −0.13695 8.35 G5 23.22 5.18 0.776 −0.181 0.10 0.0 5454 0.795 1 HD 195146 20:08:36.4 +85:07:27.2 −0.06866 −0.02057 7.13 F8 24.70 4.09 0.572 −0.046 −0.09 4.1 6038 1.916 1 HD 191806 20:09:28.3 +52:16:34.8 0.11413 0.08961 8.09 K0 13.61 3.76 0.640 −0.085 0.31 3.2 5972 2.692 6 HD 191672 20:10:21.3 +21:14:21.9 0.04189 0.01797 8.01 G2 IVws 14.44 3.81 0.544 −0.032 −0.58 2.8 5935 2.448 1 HD 193215 20:15:21.4 +64:11:55.6 0.18092 0.23031 8.44 G5 14.79 4.29 0.807 −0.206 0.10 2.7 5362 1.847 8 HD 192804 20:15:42.5 +31:14:31.7 −0.07773 −0.05442 7.68 F8 V 13.62 3.35 0.580 −0.050 −0.13 5.7 5994 3.802 1 continued on next page Appendix A. Sample stars N ) ! L (L eff (K) T i continued on next page sin v (km/s) 0.42 0.9 5772 1.482 1 0.39 3.4 6160 2.638 1 0.35 4.1 6191 2.233 1 0.12 5.2 6068 3.073 1 0.110.16 4.9 0.0 6181 1.914 5399 0.553 1 1 1.590.37 0.90.02 0.9 3.90.81 5610 1.141 1.139 5937 1.5 1 5.739 1 1 1.970 1 0.88 4.7 2.310 1 0.25 2.5 5778 3.493 1 0.15 4.6 6244 2.455 1 0.26 0.0 5337 0.531 1 0.25 2.6 4405 0.139 10 0.02 5.9 5432 0.583 1 0.13 2.2 4897 0.300 1 0.58 2.8 5907 4.040 1 0.65 3.7 6060 2.113 1 0.10 1.6 5045 0.327 2 0.18 1.9 5800 1.117 1 0.12 5.4 5875 2.782 1 0.08 2.8 6345 2.202 1 0.17 3.4 6058 2.578 1 0.29 3.8 6186 1.637 1 − − − − − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.087 0.05 3.1 5852 1.848 4 0.208 0.09 0.0 5349 0.613 1 0.067 0.013 0.012 0.039 0.100 0.24 2.90.008 5867 0.05 3.802 7.5 4 6388 2.224 3 0.025 0.167 0.101 0.046 0.023 0.067 0.009 0.080 0.130.200 5.1 0.28 5925 1.4 4.920 5457 4 1.259 2 0.078 0.184 0.050.161 0.0 0.07 5422 5.9 0.651 1 5521 0.643 1 0.100 0.15 2.0 5830 1.754 2 0.118 0.19 2.8 5753 0.716 3 0.015 0.173 0.073 0.03 2.0 5924 1.532 1 0.516 0.352 0.06 3.6 4894 0.344 3 0.130 0.27 2.2 5735 2.291 3 0.174 0.130 0.24 1.5 5726 1.500 3 0.159 0.34 0.9 5642 2.397 5 0.321 0.036 0.012 0.095 0.26 2.1 5898 2.924 2 0.276 0.080 0.069 0.14 4.0 5992 2.419 3 0.071 0.007 0.065 0.07 3.4 5987 1.288 1 0.038 0.015 0.095 0.070.104 2.9 0.12 5820 2.0 1.897 1 5795 1.681 6 0.083 0.04 2.7 5871 1.532 1 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .51 3.89 0.499 .28 4.70 0.800 .95 4.24 0.663 .45 5.61 0.766 .65 6.26 0.970 .56 3.98 0.708 .61 5.51 0.767 .98 4.44 0.707 .60 4.37 0.637 π 0.63 3.40 0.663 3.60 5.56 0.759 4.48 5.39 0.750 (mas) type Spectral V δ µ (as/yr) 0.09967 8.41 G5 14.16 4.17 0.644 0.13785 7.78 K2 34.76 5.49 0.810 0.33252 8.31 G2 V 16.43 4.39 0.610 0.14302 7.62 G0 16.50 3.71 0.503 0.00483 8.45 G0 12.23 3.89 0.490 0.00059 8.08 F8 12.56 3.57 0.560 0.09655 6.72 F8 V 29.56 4.07 0.529 0.01619 8.18 G50.24666 8.49 20.27 G5 4.71 0.665 13.02 4.06 0.572 0.89884 7.37 F8 V-VI0.07871 7.70 28.26 4.63 G5 0.525 11.09 2.92 0.610 0.02902 8.43 G0 12.150.14403 7.81 3.85 0.493 F8 11.45 3.10 0.632 0.29799 7.83 G8 IV 32.68 5.40 0.780 0.01124 8.01 G5 27.86 5.23 0.692 0.12795 8.17 G0 17.27 4.36 0.620 0.56347 8.27 K5 67.38 7.41 1.119 0.09521 7.78 G5 IV-V 17.19 3.96 0.747 0.26782 8.25 K2 42.23 6.38 0.939 0.03179 7.60 F9 Vm 18.63 3.95 0.501 0.07824 7.55 G5 16.84 3.68 0.656 0.35220 8.26 K0 39.38 6.24 0.889 0.00425 8.27 G5 19.430.18836 4.71 7.59 0.631 G2 V 17.96 3.86 0.613 0.14211 7.95 F8 14.18 3.71 0.616 0.05646 7.42 F8 19.80 3.90 0.487 0.07287 8.27 G5 17.91 4.54 0.606 0.08060 8.27 F8 12.51 3.76 0.557 0.22147 6.81 G0 30.46 4.23 0.507 0.22179 8.24 G5 15.24 4.15 0.656 0.03737 8.18 G5 V 16.70 4.29 0.670 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.27888 0.10887 0.01606 0.07974 0.12158 0.07603 0.03343 0.01183 0.02759 8.26 G0 11.00 3.47 0.629 0.23921 0.00047 0.02066 6.96 F8 V 23.24 3.79 0.507 0.06230 0.16109 0.02393 0.16365 7.75 G0 12.71 3.27 0.552 0.08138 0.01210 0.05603 0.15590 0.12465 0.09554 0.07919 0.06256 0.10072 − − − − − − − − − − − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 205351 21:33:54.3 +32:52:09.1 0.06203 HD 205286 21:33:08.8 +40:49:27.5 0.13815 HD 205027 21:32:28.5 +01:00:44.3 HD 204559HD 204712 21:28:43.4 +26:34:11.3 21:30:08.4 +12:16:14.4 0.06569 0.07696 7.80 F8 16 HD 201876 21:12:07.6 +10:38:16.8 0.17218 HD 204906 21:30:18.1 +47:10:19.0 HD 204814 21:29:46.7 +45:53:40.2 0.42657 0.35204 7.93 G8 V 3 HD 202072 21:13:22.4 +11:12:41.5HD 204155 0.16224 21:26:42.9HD 204277 +05:26:29.9 21:27:06.7 0.16709 +16:07:27.0 HD 201891 21:11:59.0 +17:43:39.9 HD 203284 21:19:36.0 +49:13:15.6 HD 205521 21:26:58.0 +81:44:32.6 0.21325 0.11966 8.06 G5 21 HD 203698 21:23:09.2 +29:38:21.5 0.00784 HD 201794 21:09:52.0 +53:44:32.2 HD 201446 21:09:30.5 +01:34:06.3 HD 203030HD 203473 21:18:58.2 +26:13:50.0 21:22:18.9BD +394559 +05:01:24.9 0.13288 21:26:56.4 0.18397 +39:39:31.5 0.00920 8.45 0.00130 0.06761 8.23 G8 V 0.01263 8.27 G5 2 G0 15 1 HD 201924 21:11:11.0 +45:27:21.3 HD 201219 21:07:56.5 +07:25:58.6 0.19167 HD 201545 21:09:45.9 +19:12:38.7 HD 201651 21:06:56.3 +69:40:28.6 0.10895 0.06605 8.19 K0 30 HD 200962 21:06:25.9 +13:06:51.3 0.03084 HD 200779 21:05:19.7 +07:04:09.5 0.07769 HD 200560 21:02:40.6 +45:53:05.1 0.40230 0.14172 7.69 K2 51 HD 200254 21:01:51.0 +21:06:19.6 0.11166 0.03049 7.64 G5 18 HD 199019 20:49:29.2 +71:46:29.3 0.13883 0.10066 8.23 G5 28 HD 200078 21:00:43.6 +17:26:52.7 0.26445 0.00460 8.05 G5 18 HD 199100 20:53:41.8 +36:07:48.1 HD 198425 20:49:16.2 +32:17:05.3 HD 197666 20:43:02.0 +52:17:43.7 HD 198109 20:47:38.4 +16:14:18.4 0.12966 HD 197623 20:44:57.1 +00:17:31.0 0.13397 HD 197396 20:42:49.4 +20:50:40.6 continued from previous page Star ID HD 193554 20:20:03.7 +23:38:17.2 HD 197488 20:42:20.4 +45:49:24.5 HD 195872 20:31:12.4 +56:53:24.3 HD 196218 20:35:42.8 +03:18:10.3 HD 197195 20:41:53.3 +12:58:49.6 BD +034368 20:34:03.3 +04:29:22.1 HD 197037HD 197207 20:39:33.0 +42:14:54.8 20:41:18.2 +30:11:28.6 0.23111 HD 196361 20:35:38.6 +36:28:29.7 HD 197140 20:40:20.9 +38:46:49.1 0.15445 0.03764 8.02 G5 18

160 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 205353 21:34:23.3 +13:55:11.5 −0.09604 −0.04999 8.18 G0 16.12 4.22 0.602 −0.062 0.06 3.0 5997 1.725 1 HD 205354 21:34:23.6 +12:15:24.5 −0.06903 −0.13888 8.05 G0 15.78 4.04 0.528 −0.024 −0.11 3.0 6185 1.966 1 HD 205702 21:36:53.3 +05:48:53.5 −0.09010 −0.11783 7.62 F8 17.46 3.83 0.560 −0.039 0.12 4.5 6167 2.419 7 HD 207897 21:40:42.9 +84:20:00.5 0.34454 0.06093 8.37 K0 34.61 6.07 0.868 −0.257 −0.11 1.2 5100 0.376 2 HD 206772 21:40:57.5 +70:32:21.9 0.14449 0.08744 8.33 G5 18.82 4.70 0.653 −0.093 0.11 2.3 5847 1.141 6 BD +374405 21:42:57.1 +38:11:04.2 −0.06883 −0.05720 8.15 G0 10.15 3.18 0.624 −0.075 −0.82 1.7 4.550 2 HD 206658 21:43:18.4 +12:31:06.0 0.07397 −0.07504 7.85 G0 20.97 4.46 0.654 −0.094 0.08 2.4 5831 1.424 1 HD 206828 21:44:09.3 +26:31:09.5 0.20345 0.05540 8.40 G2 V 17.70 4.64 0.690 −0.117 −0.03 1.3 5669 1.233 1 HD 207485 21:45:52.5 +70:20:53.0 0.12047 0.08063 8.00 G5 26.33 5.10 0.727 −0.144 −0.01 5.3 5560 0.827 1 HD 207740 21:50:36.3 +28:46:01.8 −0.04955 −0.01943 7.97 G5 V 20.27 4.50 0.737 −0.151 −0.04 1.5 5516 1.447 1 HD 207839 21:50:37.1 +49:26:16.1 −0.15919 −0.01614 7.81 K0 29.31 5.15 0.777 −0.182 0.20 1.9 5493 0.818 20 HD 207966A 21:51:53.0 +42:20:37.9 −0.17522 −0.30557 7.85 G5 33.93 5.50 0.798 −0.198 −0.06 0.3 5323 0.601 7 HD 208038 21:53:05.4 +20:55:49.9 −0.00608 −0.10170 8.18 K0 41.71 6.28 0.937 −0.319 −0.17 2.6 4886 0.328 2 HD 208552A 21:56:40.4 +16:07:25.3 −0.02736 −0.03237 8.47 F8 10.11 3.49 0.519 −0.021 0.07 4.2 6291 3.254 9 HD 208906 21:58:40.9 +29:48:45.4 −0.36174 −0.38630 6.95 F8 V-VI 34.12 4.62 0.501 −0.012 −0.94 2.4 1.140 1 HD 208863 21:58:41.4 +22:34:19.9 0.07186 −0.00779 8.28 G5 11.13 3.51 0.658 −0.097 0.18 3.5 5859 3.426 3 HD 208880 21:59:08.9 +03:11:52.0 0.22049 0.06803 8.45 G0 28.15 5.70 0.755 −0.165 0.00 1.1 5477 0.485 2 HD 209262 22:01:54.1 +04:46:13.6 0.06908 −0.10127 8.00 G5 21.86 4.70 0.687 −0.115 −0.00 1.7 5691 1.164 2 HD 209320 22:02:01.4 +23:29:44.0 0.01962 0.01545 8.33 F8 12.31 3.78 0.563 −0.041 −0.16 4.5 6041 2.537 1 HD 209393 22:02:05.4 +44:20:35.4 0.03849 0.02900 8.96 G5 29.79 5.33 0.693 −0.119 −0.19 3.8 5594 0.654 1 HD 209347 22:02:30.8 +11:22:16.7 0.15447 0.06713 8.06 G0 21.36 4.71 0.576 −0.048 −0.55 1.6 5835 1.084 1 161 HD 209599 22:04:19.9 +20:15:56.4 0.17611 −0.08594 8.37 K0 31.38 5.85 0.816 −0.213 −0.16 0.7 5229 0.442 2 HD 209858 22:05:54.0 +27:58:01.3 0.03752 −0.19343 7.78 F8 V 18.02 4.06 0.536 −0.028 −0.18 3.2 6128 1.937 1 HD 210144 22:06:42.0 +53:07:50.3 −0.53146 −0.34114 7.80 G8 V 31.99 5.33 0.788 −0.190 0.11 0.2 5423 0.698 4 HD 211681 22:06:48.5 +85:24:33.8 0.08785 0.07608 8.10 G5 14.11 3.85 0.735 −0.149 0.26 2.4 5646 2.628 12 HD 210323 22:08:45.8 +39:25:29.0 0.11874 0.05951 8.43 G0 19.22 4.85 0.634 −0.081 −0.20 2.3 5782 0.983 1 HD 210553 22:10:50.6 +15:42:29.6 0.02537 −0.02113 8.02 F8 V 22.24 4.76 0.598 −0.060 −0.26 2.2 5879 1.047 2 HD 210631 22:11:39.3 +06:11:36.4 0.23428 0.06775 8.45 G0 13.78 4.15 0.594 −0.057 −0.32 2.8 5868 1.831 1 HD 210923 22:13:03.2 +33:36:12.1 0.00852 0.06633 7.39 G0 14.72 3.23 0.559 −0.039 −0.22 3.2 6030 4.203 1 HD 210947 22:13:50.4 +00:52:35.2 0.06001 −0.03920 7.86 F8 15.41 3.80 0.528 −0.024 2.452 0 HD 211403 22:15:15.6 +58:16:31.7 0.10172 0.01640 8.48 G0 13.22 4.09 0.532 −0.026 0.27 17.6 6327 1.881 10 HD 211275 22:15:17.7 +39:46:35.1 0.23018 0.04834 7.59 G0 13.49 3.24 0.588 −0.054 0.20 3.6 6102 4.223 5 HD 211472 22:15:54.0 +54:40:22.4 0.21347 0.06991 7.50 K1 V 46.62 5.84 0.810 −0.208 0.07 3.9 5341 0.444 1 HD 211476 22:17:15.1 +12:53:54.6 0.85766 0.09818 7.04 G2 V 32.55 4.60 0.606 −0.065 −0.13 2.4 5905 1.219 1 BD +483666 22:17:26.9 +49:11:44.8 −0.11645 0.07735 8.29 F8 11.22 3.54 0.521 −0.021 −0.56 3.3 6025 3.107 1 HD 211662 22:17:44.8 +42:46:57.8 0.17151 0.07005 8.04 K0 14.26 3.81 0.582 −0.051 −0.29 2.1 5922 2.491 1 HD 211786 22:19:25.1 +12:27:35.2 0.06656 −0.09829 7.98 G5 23.88 4.87 0.666 −0.102 0.984 0 HD 212029 22:20:23.9 +46:25:05.7 −0.19034 −0.06441 8.50 G0 17.39 4.70 0.586 −0.053 −1.64 0.8 1.100 1 HD 211985 22:21:00.7 +06:35:27.3 0.10893 0.00298 7.88 F8 18.41 4.21 0.582 −0.051 −0.05 4.1 6020 1.723 1 HD 212291 22:23:09.2 +09:27:39.9 0.30593 0.03508 7.91 G5 30.84 5.36 0.682 −0.112 0.632 0 HD 212585 22:24:37.8 +39:25:50.7 0.19238 0.09242 8.03 K0 21.88 4.73 0.678 −0.110 0.09 3.5 5757 1.127 14 HD 212733 22:25:55.0 +35:21:53.5 0.18488 −0.19382 8.32 K2 34.28 6.00 0.902 −0.287 0.06 0.1 5075 0.412 8 HD 212858 22:26:43.3 +36:46:11.7 0.22656 0.10880 8.10 F8 18.12 4.39 0.584 −0.052 −0.55 1.6 5808 1.462 1 HD 213101 22:28:43.4 +26:43:32.3 0.04259 −0.00682 7.98 G0 12.81 3.52 0.611 −0.068 −0.23 2.2 5847 3.305 1 HD 213338 22:29:34.7 +52:23:24.9 0.16625 0.04583 8.41 G8 V 23.23 5.24 0.716 −0.136 0.08 1.9 5631 0.722 2 HD 213472 22:30:50.0 +40:43:19.5 −0.21458 0.04102 8.17 G5 15.54 4.13 0.700 −0.124 0.06 1.2 5674 1.984 1 continued on next page Appendix A. Sample stars N ) ! L (L 0.580 0 2.425 0 eff (K) T i continued on next page sin v (km/s) 0.04 2.5 5898 1.790 1 0.04 1.6 5194 0.376 1 0.40 2.9 5897 1.430 1 0.10 5.8 5830 1.148 1 0.13 5.1 6134 1.644 1 0.36 3.1 6136 2.601 1 0.09 2.20.22 4618 3.3 0.298 5 58990.19 1.921 5.7 1 6206 1.511 1 0.05 2.4 5874 2.457 1 0.03 2.2 5618 3.594 1 0.130.44 2.8 5.7 5238 0.639 6253 2.448 8 1 0.280.35 0.20.03 1.6 5605 5.4 2.831 5737 1.470 1 5862 4.831 1 5 0.35 2.30.21 5963 3.254 1.8 1 5660 4.341 1 0.59 1.2 5760 1.900 1 0.54 4.1 6113 2.074 1 0.20 4.9 5990 5.681 1 0.350.28 2.8 6.5 4809 0.223 6069 1.754 1 1 0.02 0.0 5162 0.5010.11 3 0.0 5509 2.526 2 0.06 2.9 6016 1.171 1 0.08 1.5 5053 0.312 2 0.880.19 2.7 4.4 5925 1.567 5.480 1 1 0.14 0.0 5708 1.127 2 0.150.19 0.0 0.0 5527 0.569 5070 0.389 1 1 0.08 5.1 6219 1.694 1 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − [Fe/H] 0.072 0.237 0.030 0.048 0.080 0.018 0.059 0.018 0.444 0.139 0.121 0.04 2.6 5678 3.314 1 0.076 0.129 0.044 0.01 3.5 60900.102 2.003 1 0.21 2.8 5845 3.844 16 0.012 0.214 0.002 0.127 0.08 3.2 5669 2.142 2 0.110 0.078 0.080 0.118 0.09 2.40.178 5712 0.20 3.156 0.8 1 5505 2.775 10 0.104 0.041 0.057 0.012 0.046 0.094 0.18 2.6 5872 3.480 3 0.320 0.030 0.249 0.146 0.136 0.10 0.8 5643 1.043 1 0.018 0.057 0.051 0.276 0.053 0.08 3.8 6060 1.876 1 0.126 0.12 4.00.100 5689 1.144 2 0.066 0.10 4.8 5990 1.550 3 0.138 0.256 0.022 BC − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − ) V − B ( V M .49 4.68 0.632 .82 5.48 0.720 .08 3.85 0.625 .12 4.41 0.629 .88 4.05 0.704 .84 3.62 0.692 .59 3.97 0.499 .10 6.70.07 4.17 0.938 0.541 .66 4.28.36 2.96 0.515 0.593 .35 4.12 0.586 .47 4.34 0.609 π 2.27 4.73 0.703 20.74 3.39 0.666 18.66 3.82 0.773 13.50 2.91 0.573 (mas) type Spectral V δ µ (as/yr) 0.20793 7.76 G0 19.28 4.19 0.619 0.28135 7.90 G5 42.63 6.05 0.845 0.07878 8.27 F8 15.63 4.24 0.540 0.05932 8.350.15615 7.86 F8 F8 11.89 3.73 20.45 4.41 0.513 0.576 0.02359 7.84 G5 19.77 4.32 0.513 0.07139 8.12 K2 47.66 6.51 1.056 0.09681 8.45 G5 10.58 3.57 0.696 0.04500 8.890.01073 G0 8.15 G0 17.00 4.04 11.72 0.569 3.49 0.706 0.03259 8.49 F8 11.52 3.80 0.500 0.10396 8.460.16172 7.07 K0 F8 25.00 5.45 21.94 3.78 0.817 0.472 0.20706 8.26 G40.01858 V 8.06 12.43 G0 3.73 0.678 10.26 3.12 0.631 0.05042 8.15 G0 11.81 3.51 0.563 0.07381 7.96 F8 V 16.97 4.11 0.593 0.07119 8.34 K2 30.30 5.75 0.859 0.07093 8.13 G5 14.21 3.89 0.730 0.04403 8.30 G5 18.42 4.63 0.582 0.20899 8.38 G5 19.58 4.840.13151 0.715 8.34 K0 38.98 6.29 0.889 0.13729 8.04 G5 21.65 4.72 0.664 0.29783 7.50 G8 V0.03423 8.32 39.82 5.50 F8 V 0.719 14.97 4.20 0.522 − − − − − − − − − − − − − − − − − − − − − − − − − α µ (as/yr) 0.19669 0.00113 0.08945 0.14063 7.98 F8 16.73 4.10 0.597 0.08381 0.02822 0.19075 0.07452 0.02766 7.95 G4 V 11.51 3.26 0.670 0.00293 0.00264 8.04 G0 IV 12.33 3.49 0.654 0.05848 0.152200.15701 0.11676 7.53 K2 50.15 6.03 0.867 0.06446 − − − − − − − − − − − − δ (d:m:s) α (h:m:s) HD 220255 23:22:13.2 +19:24:00.4 0.04466 HD 218168 23:04:02.1 +74:28:38.7 0.10738 0.03399 8.12 G5 20 HD 221851 23:35:25.6 +31:09:40.7 HD 214435 22:37:52.3 +14:59:11.3 0.03635 HD 220498HD 220842 23:23:40.3 +52:42:25.4 23:26:37.0 +56:53:11.5 HD 221822 23:35:23.2 +02:13:31.2 0.11224 0.31721 8.42 G5 25 HD 214089 22:35:18.8 +32:49:17.3 HD 218355 23:06:55.5 +38:00:59.0 0.09695HD 220221 0.02074 23:21:44.4 8.27 +45:10:33.8 G5 0.19705 13 HD 221992 23:36:31.5 +45:30:54.0 0.02311 HD 221131 23:28:25.6 +71:06:57.8HD 221573HD 0.00991 221627 23:32:53.6 +47:00:44.2 23:33:34.7 +17:49:07.8 0.25589 0.02428 6.81 G2 IV HD 220721 23:25:45.7 +45:20:50.7 0.04745 continued from previous page Star ID HD 213582 22:32:19.6 +12:14:25.0HD 214560 0.06499 22:38:54.0 +08:22:27.0 0.25774 0.02463 8.12 G0 18 HD 218059HD 218354 23:05:00.2 +04:13:31.5 23:06:53.6 +38:14:36.9 0.15393 0.07951 0.03178 8.19 G5 14 HD 214059 22:35:38.1 +05:22:24.2HD 214823 0.42672 22:40:19.8 +31:47:15.3 0.08141 HD 216320 22:50:59.1 +54:08:32.5 HD 220845 23:26:49.6 +36:06:13.8HD 0.08382 221585 0.01880 23:32:53.8 8.44 +63:09:19.8 0.43978 G5 0.03099 7.47 10 G8 IV HD 220008 23:20:18.3 +06:52:20.6 HD 218637 23:09:35.1 +01:55:34.4 0.13039 HD 221167HD 221477 23:28:49.3 +72:13:48.6 23:32:22.6 +35:19:49.7 0.05756 0.08788 0.08525 7.74 G0 17 HD 218172 23:05:35.3 +20:14:27.4 0.00711 0.02005 7.26 F8 IV HD 219828 23:18:46.7 +18:38:44.6 HD 214683HD 215151 22:39:50.8 +04:06:58.0 22:43:01.7 +14:45:33.5 0.17899 0.04578 0.10831 8.48 0.01231 8.14 K2 F8 44 16 HD 216191 22:50:14.2 +36:42:14.8 0.06807 BD +024630 23:12:36.3 +03:37:04.7 0.09364 HD 215257HD 215442 22:43:50.7 +03:53:12.6HD 215859 22:44:19.5 +51:26:58.6 0.15064 22:48:06.7HD 216175 +13:00:32.4 0.12876 0.33161 22:49:57.1 7.41 0.05363 0.07461 +50:01:00.1 7.50 F8 0.14602 F8 0.04685 23 7.92 12 G5 17 HD 221239 23:30:13.7 +31:42:20.9 0.19052 HD 219781 23:18:23.5 +12:06:48.5 HD 215274 22:43:40.4 +30:05:33.0 0.23973 0.02350 7.99 G5 V 2 HD 219428 23:15:13.3 +52:21:06.9 0.17858 0.01674 8.26 G0 16 HD 215500HD 216520 22:44:05.8 +64:34:14.4HD 215942 22:47:32.6 +83:41:49.3 0.05054 22:48:14.0 +41:31:57.0 HD 215944 22:48:29.8 +28:07:18.4

162 continued from previous page

α δ µα µδ Spectral π v sin i Teff L Star ID V MV (B − V ) BC [Fe/H] N (h:m:s) (d:m:s) (as/yr) (as/yr) type (mas) (km/s) (K) (L!)

HD 222794 23:43:26.8 +58:04:49.1 0.39005 0.48119 7.14 G2 V 22.01 3.85 0.645 −0.088 −0.71 0.0 2.484 1 HD 223061 23:46:18.4 +09:47:06.3 −0.05155 −0.18374 7.48 G0 22.20 4.21 0.582 −0.051 −0.29 1.2 5922 1.723 1 HD 223276A 23:48:08.4 +41:06:28.6 0.01997 −0.01907 8.27 G5 16.77 4.39 0.680 −0.111 −0.21 0.4 5627 1.543 1 HD 223374 23:49:01.1 +03:10:52.2 0.06447 −0.00183 8.38 G5 37.50 6.25 0.929 −0.312 −0.21 1.5 4891 0.335 1 HD 223583 23:50:44.1 +41:49:39.8 0.18087 0.06861 8.27 G0 13.76 3.96 0.578 −0.049 −0.61 2.2 5804 2.166 1 HD 223649 23:51:13.9 +62:35:33.5 0.06967 −0.09364 8.34 F8 13.03 3.91 0.569 −0.044 −0.01 4.2 6082 2.257 1 HD 223789 23:52:35.1 +35:20:11.1 0.05534 −0.05185 8.03 G5 14.69 3.87 0.548 −0.034 0.13 5.4 6213 2.321 2 HD 223848 23:52:56.2 +36:57:15.3 −0.12414 −0.09029 6.52 G0 20.07 3.03 0.571 −0.045 −0.22 3.8 5988 5.082 1 HD 224233 23:56:11.3 +59:46:03.4 0.19530 0.28218 7.67 G0 20.01 4.18 0.638 −0.084 −0.05 0.8 5831 1.826 1 HD 224508 23:58:28.0 +24:23:42.8 0.02254 −0.11030 8.38 G5 16.28 4.44 0.624 −0.075 0.09 2.7 5935 1.426 19 HD 224540 23:58:38.9 +77:15:43.2 0.01887 −0.04404 8.38 G0 23.37 5.22 0.651 −0.092 −0.29 1.8 5689 0.706 1 HD 224531 23:58:45.8 +31:56:22.2 0.07666 −0.13459 8.28 G5 20.33 4.82 0.734 −0.149 0.03 1.5 5554 1.075 1 BD +405199 23:59:06.8 +41:10:14.0 0.08220 0.00356 8.16 G0 22.91 4.96 0.652 −0.092 0.897 0 HD 224602 23:59:09.0 +41:12:06.2 0.08021 0.00440 7.72 G0 22.19 4.45 0.563 −0.041 −0.58 1.0 5868 1.369 1 163 Appendix A. Sample stars

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