Thesis
Reference
Probing the giant exoplanet and brown dwarf distribution on wide orbits with combined direct imaging and radial velocity
HAGELBERG, Janis
Abstract
Thousands of exoplanets have been discovered in the past 20 years using the radial velocity and transit techniques. But these techniques are at present limited to the detection of companions within a few astronomical units. This thesis aims at probing the giant planet and brown dwarf distribution on wider orbits. This goal is reached by combining the radial velocity and direct imaging techniques. An new data reduction pipeline called GRAPHIC has been developed in order to use the direct imaging technique. This pipeline was then used to analyse the observations taken with VLT/NACO of target who had a radial velocity trend. A statistical analysis of the observed sample was then used to derive the distribution of brown dwarfs on wide orbits. The thesis ends with the presentation of preparatory work for the oncoming observations with the new instrument SPHERE.
HAGELBERG, Janis. Probing the giant exoplanet and brown dwarf distribution on wide orbits with combined direct imaging and radial velocity . Thèse de doctorat : Univ.
Genève, 2014, no. Sc. 4698
URN : urn:nbn:ch:unige-482403
DOI : 10.13097/archive-ouverte/unige:48240
Available at:
http://archive-ouverte.unige.ch/unige:48240
Disclaimer: layout of this document may differ from the published version.
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UNIVERSITÉ DE GENÈVE Département d’Astromonie
FACULTÉ DES SCIENCES Pr. Stéphane UDRY Dr. Damien SÉGRANSAN
Probing the Giant Exoplanet and Brown Dwarf
Distribution on Wide Orbits with Combined Direct Imaging
and Radial Velocity
THÈSE
présentée à la Faculté des Sciences de l’Université de Genève pour obtenir le grade de docteur és sciences,
mention Astronomie et Astrophysique
par
Janis HAGELBERG de
Meyrin (GE)
Thèse N◦: 4698
GENÈVE
Observatoire Astronomique de l’Université de Genève 2014
The jury that assessed this thesis was composed of
Stéphane UDRY
Université de Genève, Geneva, Switzerland
Damien SÉGRANSAN
Université de Genève, Geneva, Switzerland
François WILDI
Université de Genève, Geneva, Switzerland
Yann ALIBERT
Physikalisches Institut, Weltraumforschung und Planetologie, Bern, Switzerland
Hans Martin SCHMID
Eidgenössische Technische Hochschule Zürich, Zurich, Switzerland
An Ronald...
R ÉSUMÉ
En 1995, deux découvertes majeures ont ouvert la porte au développement de nouveaux domaines de l’astronomie. D’un côté la découverte de la première planète gravitant au- tour d’un autre étoile de type solaire, et de l’autre la détection confirmée de trois objets substellaire connus sous le nom de naines brunes. Alors que les planètes étaient con- nues depuis l’antiquité grâce aux objets présent dans notre propre système solaire, les naines brunes sont des objets connus depuis peu de l’humanité étant donné qu’elles n’avaient jamais été observée avant et que leur prédiction théorique ne date que des an- nées 1960. Depuis ces premières découvertes plus d’un milliers de planètes et plusieurs milliers de naines brunes ont été découvertes. Ces nombreuses découvertes et la grande diversité d’objets ont amené avec eux leurs lots de nouvelles questions.
Ce travail de thèse s’intéresse plus particulièrement à deux types d’objets intrin- sèquement distincts qui ont cependant un grand nombre de points en commun. Ces deux catégories sont les planètes géantes et les compagnons naines brunes à longue période. En étudiant la distribution et éventuellement la photométrie de ces objets, ce travail essaie de répondre aux questions qui touchent la formation, l’évolution, et la structure de ces compagnons massif à longue période.
La méthode des vitesse radiales, utilisée pour la détection de la première exoplanète est encore à ce jour une des méthodes les plus efficaces. Le principe de la détection en vitesses radiales est de mesurer le déplacement de l’étoile hôte dans la direction de la ligne de vue, causé par l’attraction gravitationnelle de la planète qui orbite autour d’elle.
Pour mesurer ce déplacement on utilise l’effet Doppler qui déplace le spectre émis par l’étoile vers le rouge quand l’étoile s’éloigne et vers le bleu quand elle s’approche. Il n’est cependant pas possible de recourir à la méthode des vitesses radiales pour détecter de tels objets à longue période, dû à la limitation intrinsèque imposée par le durée totale des observations. Il est en effet difficile de détecter des compagnons avec des périodes plus longues que le temps pendant lequel l’étoile à été observée. De si longues périodes créent des dérives en vitesse radiale qui indiquent la présence d’un compagnon sans pour autant permettre de déterminer sa vraie masse.
L’imagerie directe a permis depuis 2008 de détecter un nombre croissant de com- pagnons sur des périodes de plusieurs dizaines voir centaines d’années, et c’est cette méthode qui est aussi utilisée dans cette thèse. Pour repousser encore plus loin les lim- ites des cette méthode, une nouvelle technique d’analyse de données a été développée dans le cadre de cette thèse, qui se base sur les transformées de Fourier. Un compte- rendus de conférence référé ainsi qu’une publication en préparation présentant cet outil nommé GRAPHIC sont inclus dans ce manuscrit.
En utilisant l’imagerie directe pour observer les étoiles à dérive en vitesse radiale il est alors possible de détecter ces objets s’ils sont assez massifs et éloignés de leur étoiles hôte. Dans le cas où il n’y a pas de détection on peut tout de même mieux cerner la nature du compagnon grâce aux limites supérieures que donne l’imagerie directe. C’est ainsi que nous avons observé 36 étoiles qui présentent un dérive en vitesse radiale. Mal- gré le fait qu’aucune détection n’ai pour l’instant pu être confirmée, ce résultat nous à cependant permis de mieux contraindre la population de planètes géantes et naines brunes à longue période. Le résultat statistique principal de ce travail est la détermi-
ii
iii
nation de la fraction de naines brunes autour d’étoiles de type solaire. Ainsi, basé sur un échantillon de 1715 étoiles nous obtenons une proportion de 0.43±0.29% compan- ion naines brunes de masse 25 à 75MJupqui orbitent autour d’étoiles FGK à un distance orbitale de 5 à 50 unités astronomiques.
Cette thèse inclut également le travail qui a été fait en préparation des observations avec le futur instrument SPHERE. Cet instrument dédié à la détection de planètes extra- solaires et qui est en cours d’installation sur le VLT permettra d’attendre des contrastes et résolution qui n’ont jamais encore été atteintes. Il sera ainsi possible d’étudier la présence de compagnons dans des domaines de masse et de séparation à ce jour in- accessibles. Le travail de préparation aux observations a amené à la publication de trois papiers acceptés pour la publication qui sont discutés et présentés dans cette thèse.
Finalement l’application de notre nouvelle réduction de données à un disque proto- planétaire a permis d’obtenir des images de la structure externe en spirale du disque à une résolution jamais encore atteinte alors. Ce résultat, basé sur des observations com- binées en infrarouge avec NICI/GEMINI et submillimétriques avec ALMA a donné lieu à la publication d’un article également inclus dans ce document.
C ONTENTS
Contents iv
1 Introduction 1
1.1 Astrophysical Context. . . 2
1.2 From clouds to stars and from disks to planets . . . 2
1.3 Planets and Brown Dwarfs . . . 3
1.3.1 How Did Giant Planets Get on Wide Orbits. . . 4
1.3.2 Brown Dwarfs Companions . . . 5
1.4 The Brown Dwarf and Giant Planet Atmospheres. . . 8
2 Detecting Sub-Stellar Companions 10 2.1 Indirect Detection Techniques. . . 10
2.1.1 Doppler Radial Velocity . . . 10
2.1.2 Photometric Transit . . . 11
2.1.3 Astrometric Motion . . . 13
2.1.4 Microlensing . . . 13
2.2 Direct Detection Techniques. . . 14
2.2.1 Direct Imaging . . . 15
2.2.2 Favoured Target Categories for Direct Imaging . . . 17
2.3 Combining the Techniques for Additional Characterisation . . . 19
3 A Technique to Directly Image Companions on Wide Orbits 20 3.1 Direct Imaging of Companions . . . 20
3.1.1 Adaptive Optics. . . 21
3.1.2 Coronography . . . 22
3.1.3 Simultaneous Differential Imaging (SDI) . . . 25
3.1.4 Angular Differential Imaging (ADI) . . . 26
3.2 GRAPHIC: The Geneva Reduction and Analysis Pipeline for High-contrast Imaging of planetary Companions . . . 27
3.2.1 Reminder on Fourier Analysis . . . 28
3.3 Publications on GRAPHIC . . . 29
GRAPHIC: IAUS 299 proceedings . . . 30
GRAPHIC: MNRAS draft . . . 32
3.4 Further improvements . . . 40
4 Using Radial Velocity Drifts to Select Candidate Companions for Direct Imag- ing 41 4.1 The Geneva Radial Velocity Surveys . . . 42
4.1.1 The CORAVEL Legacy . . . 42
4.1.2 The CORALIE Survey. . . 43
4.1.3 The HARPS Survey . . . 44
4.2 Radial Velocity Trends. . . 44
4.2.1 Building a coherent sample from the radial velocity surveys . . . 45 iv
CONTENTS v
4.2.2 Target sample selection . . . 45
4.2.3 Probability Maps Built on the Radial Velocity Data . . . 46
4.3 Direct Imaging Observations. . . 49
4.3.1 Observing with NaCo . . . 49
4.3.2 Reducing the data . . . 51
4.3.3 Analysing the observations . . . 51
5 Results of the Direct Imaging Campaign with NaCo 55 5.1 General notes on the presentation of the observations . . . 56
5.2 Confirmed radial velocity detections . . . 58
5.2.1 HD142 . . . 58
5.2.2 HD361 . . . 58
5.2.3 HD68475. . . 59
5.2.4 HD78746. . . 60
5.2.5 HD98649. . . 63
5.2.6 HD135625 . . . 65
5.2.7 HD145666 . . . 65
5.3 Companion Candidates . . . 71
5.3.1 HD139189 . . . 71
5.3.2 HD168863 . . . 72
5.3.3 HD170209 . . . 72
5.3.4 HD178340 . . . 73
5.4 Non-detections . . . 78
5.4.1 HD4113 . . . 78
5.4.2 HD11397. . . 78
5.4.3 HD101339 . . . 79
5.4.4 HD117939 . . . 79
5.4.5 HD124553 . . . 80
5.4.6 HD139879 . . . 80
5.5 Visual binaries . . . 80
5.5.1 HD93351. . . 86
5.5.2 HD104263 . . . 87
5.5.3 HD111031 . . . 87
5.5.4 HD150139 . . . 88
5.5.5 HD185615 . . . 88
5.5.6 HD197823 . . . 89
5.5.7 HD207700 . . . 89
5.5.8 HD212708 . . . 90
5.5.9 HD213941 . . . 90
5.5.10 HIP948 . . . 91
5.6 Additional targets from the previous CORALIE-NaCo+SDI survey . . . 91
6 Constraints on Wide Orbit Companions from Combined RV and Direct Imag- ing 94 6.1 Overall Survey . . . 94
6.1.1 Binary rate. . . 95
6.2 Brown Dwarf Occurrence Rate. . . 97
6.2.1 Result from this survey . . . 97
6.2.2 Occurrence rate from the previous CORALIE-NaCo+SDI survey . 98 6.2.3 Occurrence rate from other surveys . . . 99
6.3 Conclusion . . . 102
7 Closer and Smaller Companions with SPHERE 103 7.1 Scientific Goals. . . 104
vi CONTENTS
7.2 The Instrument . . . 105
7.2.1 The Common Path and Infrastructure (CPI) . . . 105
7.2.2 The InfraRed Dual Imaging Spectrograph (IRDIS). . . 106
7.2.3 The Integral Field Spectrograph (IFS) . . . 109
7.2.4 The Zurich Imaging Polarimeter (ZIMPOL) . . . 109
8 Getting Ready for SPHERE 113 8.1 The NIRSUR Target Sample Preparation . . . 114
8.2 The NaCo Large Program (NaCo-LP) . . . 115
8.2.1 I - Sample definition and characterization . . . 117
8.2.2 II - Survey description, results and performances. . . 135
8.2.3 Astrophysical false positives in direct imaging for exoplanets: a white dwarf close to a rejuvenated star . . . 154
8.3 Planets in Reflected Light. . . 167
9 GRAPHIC on Extended Sources 170 9.1 Angular Differential Imaging on Extended Sources. . . 170
9.2 Direct Detection of Gas Flows through a Protoplanetary Gap. . . 170
10 Conclusion & Perspectives 184
A Non Drifts 186
Bibliography 191
Thanks 211
C H A P T E R
1
I NTRODUCTION
Chapter contents
1.1 Astrophysical Context. . . 2
1.2 From clouds to stars and from disks to planets . . . 2
1.3 Planets and Brown Dwarfs . . . 3
1.4 The Brown Dwarf and Giant Planet Atmospheres. . . 8 The present work focuses on two different type of objects which even though they are intrinsically different share many important characteristics. These two object categories are the giant planets and the brown dwarfs companions on wide orbits. By exploring their distribution on wide orbits and possibly probing their atmospheres this work aims at shedding some new light on various open questions about giant planet and brown dwarf formation which will be introduced in this thesis.
In this first chapter I will introduce these two categories of objects, followed by the open questions this work tries to address about the formation, distribution, and struc- ture of these objects. Current knowledge of the possible formation mechanisms strug- gles to explain the existence of the detected massive companions on wide orbits, but the population of these wide companions is not yet well understood due to detectabil- ity constraints. Finally, observations which where able to probe their atmospheres shed some light on the complex photo-chemistry of their upper layers.
These questions would not have been triggered if there had been no objects to work on, detected by a variety of methods. Thesecond chaptergives an overview of these di- verse planet detection techniques emphasising on the mass-separation range and phys- ical parameters they give access to, as well as the favoured targets they can probe and their intrinsic limitations.
This thesis mainly relies on the direct imaging technique in order to contribute to the debate on the formation, distribution, and structure of massive companions on wide orbits. The technique as well as the associated pipeline specifically developed for this project are described in thethird chapter. The first published results of this pipeline are presented in chapter9where the pipeline was used to search for companions around a star hosting a circumstellar disk. Images of the disk were also reduced by this pipeline and included in the publication.
1
2 CHAPTER 1. INTRODUCTION
The core of this work which aims at directly detecting companions causing a trend in the radial velocity measurements is presented in chapters4to6of this thesis. Chapter 4describes the target sample and the way it was built. The reduced observations of the companion candidates and non-detections is given in chapter5, while the statistical analysis of the whole observing project is given in chapter6.
The upcoming instrument SPHERE/VLT which is specifically designed for the direct detection of sub-stellar companions is described in chapter7, followed by chapter8on the preparatory work for the guaranteed time observations to which I contributed.
Thelastchapter concludes this work and presents perspectives in the near and far future. Supplementary data is presented in the appendix.
1.1 Astrophysical Context
I will start by shortly introducing the mechanisms leading to the formation of a plan- etary system emphasising the difference between stellar and planetary formation. It is followed by a description of the two type of objects the present work is focusing on and the different issues linked to them.
1.2 From clouds to stars and from disks to planets
For the sake of brevity and simplicity I will only explain the formation scenario of an isolated low-mass star, which starts with a molecular cloud becoming unstable. A more detailed treatment can be found in Shuet al.(1987) and McKee & Ostriker (2007) where more complex parameters such as multiple systems, turbulence, asymmetries, mag- netic fields, and rotation are taken into account. More details on the particular case of brown dwarfs are given in section1.3.2.
Once the cloud becomes gravitationally unstable it starts to collapse (Jeans,1902;
Hoyle,1953; Hunter,1962). At the beginning the density of the gas is still low so that the gravitational energy can be entirely radiated through the optically thin gas. The isothermal collapse stops when the gravitational heating overwhelms the radiative cool- ing which is impeded by the increased opacity. A first core is formed at this stage as the contraction is slowed down.
The contraction slowly continues nonetheless as the surrounding gas continues to in-fall on the core, until the gravitational heating reaches a temperature of about 2000K where hydrogen molecules start to dissociate. By absorbing the gravitational energy, the dissociation triggers the second collapse of the core as gas pressure cannot compensate for gravitational pressure any more (Masunagaet al.,1998). This slowing down followed by a collapse repeats again when hydrogen and then later helium get ionised.
Due to conservation of angular momentum, any slight rotation of the initial molecu- lar cloud translates into a much faster rotation of the final core and its surrounding enve- lope (Kant,1755; Laplace,1808; Hoyle,1960). As a result, the remaining envelope around the star condensates into a circumstellar disk. The first unresolved discovery of such a circumstellar disks was achieved with theInfrared Astronomical Satellite(IRAS) through the photometric detection of infrared excess of four main sequence stars (Neugebauer et al.,1984; Gillett,1986; Stromet al.,1989). Grains in the circumstellar disk absorb short wavelength emission from the star thus heating them up, these warm grains then re-emit this energy in infra-red resulting in a spectral energy distribution which can be approximated as the composite of two black bodies.
This circumstellar protoplanetary disk is the birthplace of planets, and the observed lifetimes of
∼<7M yras well as the typical masses of 10−3−10−1M¯of such disks around sun-like stars are among the strongest constraints on planetary formation models (Beck- with & Sargent,1996; Haischet al.,2001; Hernándezet al.,2007). Resolved infra-red ob- servations of evolved disks revealed the formation of spiral structures and gaps, while
1.3. PLANETS AND BROWN DWARFS 3
Figure 1.1– Left: Discovery image of Gliese 229B with ground based adaptive optics imaging on the Palomar 60-inch telescope. Right: Confirmation of discovery with second epoch ob- servation using the Hubble Space Telescope (Nakajimaet al.,1995).
recent high resolution sub-millimetric observations with theAtacama Large Millime- ter/submillimeter Array(ALMA) started to give access to dynamic structures of gas and dust particles in the disk (Acke & van den Ancker,2006; Fouchetet al.,2010; Dodson- Robinson & Salyk,2011; Ayliffeet al.,2012; Casassuset al.,2013). The detailed process of sedimentation and coagulation that forms planets starting from small dust grains and gas present in the disk remains a debated topic, specially with regards to the observed planet population (sec.1.3.1).
1.3 Planets and Brown Dwarfs
In 1995, two key discoveries triggered the rapid expansion of new research fields in as- tronomy. On one hand, the first detection of a planet orbiting a sun-like star set off the whole field of exoplanet studies (Mayor & Queloz,1995). On the other hand, the un- ambiguous detection by three independent teams of three different sub-stellar objects known as brown dwarfs (Basri & Marcy,1995; Nakajimaet al.,1995; Reboloet al.,1995, and figure1.1) settled the existence of these peculiar objects which were by that time only theoretically predicted. Unlike exoplanets which are objects already known from our own Solar system, brown dwarfs were purely theoretical objects predicted with no observation for more than 30 years before their discovery (Kumar,1962,1963).
Eighteen years after these first discoveries, more than 1000 exoplanets and nearly 2000 brown dwarfs have been discovered, and about 3000 candidates from the Kepler space mission are waiting to be confirmed. These numbers are growing ever faster as the pace of new detections is increasing, thanks to newly built instruments purposely designed to search for sub-stellar objects, but also thanks to the optimisation of data analysis techniques and observation strategies.
4 CHAPTER 1. INTRODUCTION
1.3.1 How Did Giant Planets Get on Wide Orbits
The discovery of 51 Peg b, which turned out to be a jupiter mass planet on an orbit ex- tremely close to its host star (Mayor & Queloz,1995) was the first hint that the previously well accepted paradigm of planet formation had to be reconsidered. This first detection was quickly followed by an avalanche of similar hot jupiter detections (Butler & Marcy, 1996; Butleret al.,1997; Noyeset al.,1997; Butleret al.,1998; Delfosseet al.,1998; Marcy et al.,1998; D. Fischeret al.,1999) which triggered new developments in planet forma- tion and migration modelisation. More recently, the discoveries through direct imaging of giant planets on wide orbits challenged these models once more (Maroiset al.,2008;
Lagrangeet al.,2010; Maroiset al.,2010; Kuzuharaet al.,2013; Rameauet al.,2013b), as the two main scenarios for planet formation, core accretion and gravitational instability, both struggle to explain the formation of massive companions on such wide orbits.
Core accretion, which is based on dust grains accreting to form planetesimals which in turn accrete into planets, can explain the formation of small planets (Wetherill,1980;
Pollacket al.,1996; Ida & Lin,2004; Alibertet al.,2005). However, the mechanism would take longer than the lifetime of the protoplanetary disk to form planets of few jupiter masses on orbits beyond∼10 AU from a Sun-like star (Pollacket al.,1996; Ida & Lin, 2004)
On the other hand, gravitational instability, where the protoplanetary disk collapses due to its own weight could explain the formation of massive companions (Cameron, 1978; Boss,1997). Nonetheless, gravitational instability needs unusually cool and mas- sive disks which are not compatible with current disk models and observations. Alter- native formation mechanisms, mainly based on these two scenarios have since been developed to try to reproduce the observed planet distribution.
Orbital evolution after formation such as migration induced by planet-disk inter- action (Linet al.,1996), planet-planet (Kozai,1962), and even planet-star interaction could explain how planets exist at orbits where no known formation mechanism can form them. The planet-disk interaction is the most efficient migration process and there are currently three type of migration mechanisms known. Which one of these different migration mechanisms comes into play mainly depends on the planet mass and disk density.
In type I migration the planet perturbation is small enough to preserve the overall structure of the disk, nonetheless forming spiral density waves where the inner wake exerts a positive torque and the outer wake a negative torque on the planet (Goldreich
& Tremaine,1979,1980, Fig.1.2a). It also generates a perturbation on its orbital radius where disk material turns in a horseshoe shaped region (Ward,1991, Fig. 1.2b). The difference of torque applied by the inner and outer spiral waves on the planet migrate it inwards (Ward,1997). The efficiency of this migration mechanism is so high that no planet would remain on orbits of several AU, but more detailed models of the process including 3D modelling, non-isothermal disks, magnetic fields, or inclination and ec- centricity of the planet orbit led to some hampering of the migration rate (Baruteauet al.,2013, and references therein).
Type II migration occurs when the planet is massive enough to carve out an annular gap in the disk. Similar to the spiral waves of type I migration, the planet takes angu- lar momentum from the inner disk while giving angular momentum to the outer disk.
As the inner disk accretes on the star the planet which is locked in its gap moves in- ward on a disk-viscosity dependent accretion timescale (Papaloizou & Lin,1984; Lin &
Papaloizou,1986a; Lin & Papaloizou,1986b). The type of gap opening and migration depends on the planet mass, disk density and viscosity profile (Cridaet al.,2006). If the inner disk becomes much less massive than the planet, the outer disk is not any more pushing the planet inwards as fast as the inner disk accretes which leads to the inner disk depletion (Quillenet al.,2004). The outer disk then continues to accrete on the planet, pushing it inwards while the planet cannot any more take up angular momentum from
1.3. PLANETS AND BROWN DWARFS 5
the depleted inner disk. Taking into account the differential gas and dust opening Crida
& Morbidelli (2007) have shown that type II migration can also be directed outwards for large disk viscosities.
Such gaps possibly caused by a planet may have recently been detected, and if con- firmed could validate this type of migration (Debes et al., 2013; Quanzet al., 2013).
Meanwhile, cavities where the inner disk has disappeared have already been observed frequently. Theses detections based on high resolution imaging are sensitive to light scattered by specific dust grains sizes tracing only part of the disk structure (Fouchet et al.,2010). With the resolution offered by the ALMA and its ability to probe a wide range of wavelengths and thus covering various dust sizes and gas emission lines it is now possible to obtain the whole picture of gas and dust dynamics in disks. Such a disk observation with ALMA can be found in section9.2, where the co-authored article by Casassuset al.(2013) is presented.
A third type of migration happens when the migration rate is not anymore deter- mined by the disk torque but rather by the migration rate itself, leading to feedback which can lead to a runaway migration until the so-called fast regime is reached. As for the two other migration types, type III migration can also be reverted leading to out- ward migration (Masset & Papaloizou,2003). A planet can go through different migra- tion types during its early evolution, and the complexity of the process increases when considering multiple planets interacting which often lead to resonant systems.
Even though these migration mechanism can move planets outwards, they still strug- gle to reproduce the observed cold jupiter population recently uncovered through direct imaging. Migration involving two giant planets where the gap overlaps can move these planets further out but not as far as the observed cold jupiters, and they rely on very specific planet mass-ratios and disk densities (Masset & Snellgrove,2001). Scattering in two or three planet systems could place one or more giant planet on the observed wide orbits, and even be at the origin of the observed free-floating planets (Chatterjeeet al., 2008; Marzariet al.,2010; Moeckel & Armitage,2012). Similar planet-planet interac- tion for planets formed by gravitational instability seems easier as the planet is initially formed on wider orbits than through core-accretion. A detailed description of the dif- ferent migration mechanism can be found in the review article by Baruteauet al.(2013).
An alternative to planet migration through interaction with the disk is the process referred to as Kozai cycles with tidal friction (Kozai,1962; Kiselevaet al.,1998; Eggle- ton & Kiseleva-Eggleton,2001). In this three-body process the orbital plane of the in- ner companion precesses due to perturbation of the outer companion with a different orbital inclination. This instability induces eccentricity oscillations on the inner com- panion and thus tidal friction which finally translates into heating detectable through increased luminosity. Such a mechanism works on much longer time-scales than the disk lifetimes.
Better characterisation of the wide orbit companion population through direct imag- ing should further constrain these formation and migration mechanisms (sec2.2).
1.3.2 Brown Dwarfs Companions
The current definition of brown dwarfs, which was adopted in 2002 by the Working Group on Extrasolar Planets of the International Astronomical Union (Bosset al.,2003), relies solely on the capacity of such objects to burn deuterium but not hydrogen in their core. Most models show that such objects with masses above≈13MJupburn deuterium in their core through the reaction
2D+1H→3He+γ
(Saumonet al.,1996; Burrowset al.,2001, figure1.3). The lower mass limit is given by the minimal mass where deuterium fusion can happen, while the upper mass limit is given
6 CHAPTER 1. INTRODUCTION
(a)Spiral density wave (b)Horseshoe density perturbation.
Figure 1.2– Type I migration of a 5 M⊕planet. White curves and arrows in figure(b)show typical gas trajectories along the horseshoe region. (Baruteauet al.,2013)
by the minimal mass where hydrogen fusion can start (Chabrieret al.,2000; Spiegelet al.,2011). As a result, the working definition for these objects is merely based on their mass, between about 13 MJupand 75 MJupwhich bridge the mass-gap between planets and stars.
But this definition is highly debated as the upper and lower limits depend on the models used but also on the initial conditions. Furthermore, the discovery of HR8799b, c, and d with masses between 7-10 MJupon wide orbits (27-68 AU) (Maroiset al.,2008) and other the hand Corot-3b which is a≈22 MJupobject on a very close orbit (0.057 AU) around its star (Deleuilet al.,2008) further advocates in favour of a definition based on formation instead of mass (Chabrieret al.,2005). This way, objects forming in proto- planetary disks would be planets, and objects forming in proto-stellar clouds would be brown dwarfs or stars depending on their ability to sustain hydrogen fusion. In this context, the lower mass limit for brown dwarfs could be as low as the minimal mass for opacity-limited fragmentation in turbulent cloud cores (Silk,1977), which is thought to be around 3 MJup(Kumar, 2003; Boyd & Whitworth,2005). Recent detections of free floating low-mass objects come in support of such a low limit (Zapatero Osorioet al., 2002; (MOA) Collaboration & (OGLE) Collaboration,2011). Other formation mechanism may be able to form free floating objects of even smaller mass such as photo-erosion of pre-existing stellar cores (Hester,1997; Whitworth & Zinnecker,2004), or the ejection of proto-stellar cores due to interactions between discs in dense clusters (Reipurth &
Clarke,2001; Bateet al.,2002; Shenet al.,2010; Thieset al.,2010). For a given mass we can a priori find both types of formation, some who formed as planets and other who formed as stars.
The problem of such a definition is to know how the massive companions have formed in hindsight. For very young objects, the luminosity is expected to depend on the formation mechanism, but this formation imprint on the luminosity rapidly fades away as the object cools down on a Kelvin-Helmholtz time-scaleG M2/RL, which is of the order of 100Myr for 5MJupand 1Gyr for 10MJup(Chabrieret al.,2006; Marleyet al., 2007; Fortneyet al.,2008b). Following the recent direct detections of young giant ex- oplanets, models with various initial conditions were calculated to specifically test the early evolution of such objects. This was done either by using a grid of initial entropies in order to cover all the possible formation specificities (Spiegel & Burrows,2012; Mar-
1.3. PLANETS AND BROWN DWARFS 7
Figure 1.3–Evolution of the luminosity (in L¯) of isolated solar-metallicity red dwarf stars and substellar-mass objects versus age (in years). The stars are shown in blue, those brown dwarfs above 13 MJupare shown in green, and brown dwarfs/giant planets equal to or be- low 13 MJupare shown in red.Figure and caption from Burrowset al.(2001).
leau & Cumming,2014), or by modelling specific formation processes and using these results for the initial entropy (Mordasiniet al.,2012b; Mordasini,2013). Even for one specific type of formation mechanism such as the core accretion the initial entropy is highly unknown as it depends on the detailed physics of the accretion shock, leading to a wide range of luminosities.
Objects formed in a circumstellar disks are expected to have higher metallicities than their host star due to the planetesimal bombardment they endure, which in turn should enrich their interiors and atmospheres (Marleyet al.,2006). Direct imaging, with the ability to provide photometry and even spectroscopy (figure2.9) of brown dwarfs and giant planets at different evolutionary stages plays a key role for the understanding of formation and atmospheres of these objects. This spectroscopic examination is hin- dered by the important role of clouds, non-equilibrium chemistry, and photo-chemistry (Chabrieret al.,2006; Hellinget al.,2008a; Baraffeet al.,2010).
Another open question is to know how the population of sub-stellar companions is distributed on wide orbits, which is tightly linked to the previous question on the forma- tion of these objects. From the nearly 2000 confirmed brown dwarf detections most are free-floating field dwarfs, while some are part of very low mass binaries (Mtot∼<0.2MSun), and only a few tens are close (∼<1000 AU) companions to sun-like stars were the bound- ary between planet and brown dwarf becomes fuzzy (online catalogue based on Dupuy
& Liu,2011; J. Schneideret al.,2011, and DwarfArchives.org). In addition to that, ra- dial velocity surveys find that∼14% of solar-type stars have planets within 4.5 AU with Msin(i)=0.15–13 MJup(Udry & Santos,2007; Mayoret al.,2011). Within 10 AU,∼13%
have stellar companions with masses above Msin(i)=80 MJup (Duquennoy & Mayor, 1991; Halbwachset al.,2003), but for brown dwarf companions with masses Msin(i)=13–
8 CHAPTER 1. INTRODUCTION
Figure 1.4– Cumulative mass distribution of potential brown-dwarf companions in the CORALIE survey. The blue dashed line shows the distribution of all 21 candidates. The black solid line shows the distribution for the 11 remaining candidates after removal of the 10 stellar companions. The companion of HD 38529 with M2 = 17.6 MJup(Benedictet al.,2010) is included. For comparison, the red dash-dotted line shows the expected cumulative distri- bution if all companions had masses of 80 MJupbased on the assumption that the orbits are randomly oriented in space.Caption and figure from Sahlmannet al.(2010).
80 MJupon the same separation range the rate goes down to 0.6% (Sahlmannet al.,2010).
This scarcity of brown dwarfs on short to intermediate periods around solar-type stars is known as the “brown dwarf desert” (Marcy & Butler,2000; Mayor & Udry,2000; Udry
& Mayor,2001; Grether & Lineweaver,2006, figure1.4). It is not clear what the brown dwarf population on wider orbits is and what the orbital extend of this desert is, as only recent technical breakthrough in direct imaging made it possible to probe these regions.
1.4 The Brown Dwarf and Giant Planet Atmospheres
Even though brown dwarfs and giant planets may form in a different way, their atmo- spheric properties are very similar (Fahertyet al., 2013; Liuet al., 2013; Showman &
Kaspi,2013). In addition to that, most current atmospheric models do not take into consideration the formation mechanism, so that a key parameter for the models which is the initial entropy remains completely unconstrained. As a result, models are used for planet and brown dwarf observations equivalently.
The recent direct detections of massive planets have given first information about their young thermally stable atmospheres, as opposed to the spectral probes of highly irradiated transiting giant planets. Photometry of these objects on wide orbits with dy- namically determined mass has shown that current atmospheric models tend to overes- timate luminosity (Closeet al.,2007; Liuet al.,2008; Dupuyet al.,2009). Low resolution spectroscopy and time resolved photometry of massive companions gave evidence to dust cloud structures strongly varying in time (Artigauet al.,2009; Buenzliet al.,2012;
Radiganet al.,2012; Apaiet al.,2013; B. A. Billeret al.,2013a; Heinzeet al.,2013). Recent models which include improved treatment of cloud condensation to take these new ob-
1.4. THE BROWN DWARF AND GIANT PLANET ATMOSPHERES 9
servations into account seem to better fit observational results for young objects and their apparent under-luminosity (Madhusudhanet al.,2011; Deaconet al.,2012; Morley et al.,2012). The current main challenge for atmospheric models remains the dynamical treatment of clouds and the condensates composition. As a result, the variety of models such as DUSTY (Chabrieret al.,2000, figure2.10a), COND (Baraffeet al.,2003, figure 2.10a), DRIFT (Hellinget al.,2008b), BT-SETTL (Allardet al.,2013), theTucsonmodels (Burrowset al.,2006; Saumon & Marley,2008, figure2.10b), and so forth mainly differ by their treatment of cloud condensation (Leggettet al.,2013; Meshkatet al.,2013).
Well constrained atmospheric models of giant planets and brown dwarfs are crucial for the whole field of direct imaging, as most direct detections of sub-stellar field objects and companions solely rely on the atmospheric models to determine their mass. Any adaptation of the models will impact a large number of already photometrically deter- mined masses and will thus potentially change our knowledge of the planet population on wide orbits. Furthermore, recent direct imaging campaigns provided relatively young objects as they are easier to find. This means that atmospheric models could not yet be probed for long term evolution.
The main project of this thesis is to characterise the brown dwarf companion popu- lation with combined radial velocity and direct imaging, in a separation regime between the regions which can be probed by these two techniques alone. It also paves the way for SPHERE by finding targets where the companions will be within reach and that will have a mass that can be dynamically determined by existing radial velocity data (sec. 2.1.1, and chap. 4). Detecting such older objects would help constraining the models inde- pendently of the formation process, as they would have had enough time to cool down.
The brown dwarfs this work focuses on are those orbiting Sun-like stars. Throughout this work, the term brown dwarf thus refers to brown dwarf companion as opposed to field brown dwarf, if not otherwise specified.
C H A P T E R
2
D ETECTING S UB -S TELL AR C OMPANIONS
Chapter contents
2.1 Indirect Detection Techniques. . . 10 2.2 Direct Detection Techniques. . . 14 2.3 Combining the Techniques for Additional Characterisation . . . 19 The current planet detection methods can be classified in two categories, depending whether they directly detect a signal coming from the planet or whether it is the indirect effect of the planet on the parent star which is measured.
The majority of exoplanets were discovered with the radial velocity or transit tech- niques. These two methods, along with astrometry do not detect the companion itself but rather the effect it has on its host star. An important characteristic of these indirect techniques is that the system needs to be observed at least as long as the orbital period of the companion in order to have the full phase coverage. This implies that companions with increasing period get detected as the duration of such surveys increases.
2.1 Indirect Detection Techniques
Even though indirect detection seem less obvious, Doppler radial velocity and photo- metric transit are currently by far the most efficient detection methods. These two tech- niques hold for more than 90% of all planet detections. While direct imaging, microlens- ing, and astrometry account for less than 10% of the detections.
2.1.1 Doppler Radial Velocity
The radial velocity technique measures the Doppler shift of the stellar spectrum caused by the movement of the star. A star hosting a companion will orbit around the centre of mass of the star-companion system as does the companion. From the measured radial velocity variation of the star it is then possible to derive the periodP, eccentricitye, ar- gument of periastronω, and temporal offsetT0. If the massm1of the host star is known it then also possible to derive the minimum massm2sini of the companion. While the true mass of the companion is not accessible due to the unknown orbital inclinationi. The link between the different parameters can be seen in equation2.1which gives the
10
2.1. INDIRECT DETECTION TECHNIQUES 11
(a)Observations at four different epochs (b)Phase-folded observations
Figure 2.1– Doppler radial velocity data which resulted in the discovery of 51 Peg b by Mayor &
Queloz (1995), observed with ELODIE/OHP.
semi-amplitude of the host starKSassuming that its massmSis much larger than the companion massmC:
KS= 28.4 p1−e2
µm2sini MJup
¶ µmS M¯
¶−2/3µ P year
¶−1/3
[m/s] (2.1)
This technique, already routinely used for the study of binary stars, is the one used by Mayor & Queloz (1995) to detect the first exoplanet around a sun-like star (figure 2.1). Many other exoplanets were later found with this technique. Until 2013, radial velocity still had the largest share of detections with more than half of all the exoplanets detected (J. Schneideret al.,2011). Not only were single planet systems discovered by radial velocity, but also multiple planet systems with 5 to 6 planets (D. A. Fischeret al., 2008; Loviset al.,2011b), and some systems packed within 0.5 AU (Boruckiet al.,2011;
Lissaueret al.,2011).
The optimal target types for this method are bright stars with numerous strong ab- sorption lines in their spectrum, low activity, and slow rotation. It is currently hard to obtain accurate radial velocity measurements on young stars which usually have high activity and rotation, as well as on massive stars with limited spectral features. The ra- dial velocity measurements have become so accurate that stellar activity has become the major limitation even for the quitest stars (Dumusqueet al.,2011b,c; Loviset al.,2011a).
By monitoring activity indicators such as the R’HK(Wilson,1963,1968), the full width at half maximum (FWHM) of the cross-correlation function (CCF), and the CCF bisector it is possible to de-trend the radial velocity data to some extend in order to detect the sig- nal caused by the companion, such as for the case ofαCentauri B b (Dumusqueet al., 2012).
2.1.2 Photometric Transit
The transit technique uses the host star photometric flux variation caused by the com- panion when it passes in front of the stellar disk. The very specific geometric configu- ration needed for a transit to be observable from Earth results in extremely low transit probabilities. Nonetheless, one of the earliest times it was mentioned as a method for planet detection was by Struve (1952). A deeper methodology was later developed by
12 CHAPTER 2. DETECTING SUB-STELLAR COMPANIONS
(a)Time series (b)Binned data
Figure 2.2– Photometric transit detection of HD 209458 b.(a)shows the time series corrected for grey and colour dependent extinction, while(b)shows the same data binned in 5 minutes averages. (Charbonneauet al.,2000)
Rosenblatt (1971). Finally, in 2000, two independent teams detected the first transit of a planet previously known by radial velocity (Charbonneauet al.,2000; Henryet al.,2000, figure2.2).
Since only a small fraction of stars have a probability to exhibit a transit, the strat- egy adopted by most transit search surveys is to observe a large number of stars. Fur- thermore, the transit probability decreases with orbital period while the detectability is proportional to the planet radius. As a results, most planets detected from the ground by this method are large short period planets known as hot jupiters. Another important draw-back of the transit method is the high false-positive rate. This requires additional follow-up observation to ensure that the transit signal is not due to one of various con- figuration of eclipsing binaries, or to a giant planet transiting a companion star which mimics a smaller planet transit (Fressinet al.,2013). These follow-up observations are carried out using direct imaging and mostly radial velocity, which has the advantage of providing the mass. In the case of Kepler giant planet candidates, Lillo-Boxet al.(2012) found that∼42% of 98 surveyed candidates had nearby objects using lucky imaging, while Santerneet al.(2012) obtained a false-positive probability of∼35% by following up 46 Kepler candidates using radial velocity. In some cases of multiple planet sys- tems a transit timing variation can be observed which can be due to the rotation of the star around the multi-planet centre of mass making it possible to detect further non- transiting planets (Miralda-Escudé,2002; Holman & Murray,2005; Ballardet al.,2011;
Barroset al.,2014).
Even though transit detection suffers from a high false-positive rate, requiring thor- ough follow-up observations to confirm the candidates, the effort is worth the benefits as transiting planets can deliver a wealth of information specially when combined with radial velocity observations (section2.3). In favourable cases the atmospheric compo- sition of these short period highly irradiated planets can be probed with two spectro- scopic methods. When the planet transits, some light from the star goes through the planets atmosphere which can then be measured with transmission spectroscopy (e.g.
Charbonneauet al.,2002; Fortneyet al.,2008a; Deminget al.,2013, figure2.3a). Most acquired spectra of small mass planets are featureless which is thought to be due to the presence of haze and clouds. Another way to probe the transiting planet atmospheres is to observe them at occultation, when they pass behind the star, which gives access to the thermal emission of the planet where absorption features can be measured (Char- bonneauet al.,2005; Deminget al.,2005; Grillmairet al.,2008, figure2.3b).
2.1. INDIRECT DETECTION TECHNIQUES 13
(a)Transmission spectroscopy
(b)Emission spectroscopy
Figure 2.3–(a)transmission spectroscopy of HD 209458b (points and boxes) compared to differ- ent models (blue and red line) (Deminget al.,2013, and referenced therein).(b)shows the emission spectroscopy measured with the occultation technique, where as before the points stand for the observations and the lines for various models (Grillmairet al.,2008, and references therein).
2.1.3 Astrometric Motion
The astrometric motion method relies on the same physical effect as the radial velocity method, which is that a star hosting a companion will orbit around the centre of mass of the system. By precisely measuring the position of a star on the sky it should be possible to detect its movement induced by the orbiting companion. All orbital parameters of an astrometrically detected companion can be determined, including the true mass. The detectability increases with companion mass and orbital period, while it decreases with host star mass and distance. Currently no planet has been discovered using astrometry, but precisions of 40-50µas have already been achieved, enough to detect jupiter-mass companions (Benedictet al.,2002; Lazorenkoet al.,2009; Sahlmann,2012). Astrometry has been successful at determining the mass of known planets previously discovered by radial velocity (Sahlmannet al.,2010,2011).
Astrometry has a promising future with the recent successful launch of theGlobal Astrometric Interferometer for Astrophysics satellite(GAIA) (Casertanoet al.,1996; Lin- degren & Perryman,1996) which is expected to detect several thousands of giant plan- ets out to 3-4 AU and hundreds of multi-planet systems (Casertanoet al.,2008; Sozzetti et al.,2014).
Even without any discovery through astrometry yet, this technique already plays an important role in combination with other detection techniques, such as radial velocity and direct imaging (section2.3).
2.1.4 Microlensing
Light rays passing near a massive body get bent by its gravitational field. This bending phenomenon was first observed during the solar total eclipse in May 1919 (Dysonet al., 1920), following predictions by Einstein based on his theory of general relativity and thus giving a first observational proof of this theory. The lensing effect, where a star bright- ness increases due to this bending was later detailed by Einstein (1936) even though only the lensing caused by the Sun could be observed at that time due to instrumental limitations.
The detection of planets by microlensing was first considered by Liebes (1964), but the first extrasolar lensing phenomenon observed was caused by a galaxy which, due to
14 CHAPTER 2. DETECTING SUB-STELLAR COMPANIONS
(a)LRG 3-757
(b)First planet discovery by microlensing
Figure 2.4– The strong gravitational lens LRG 3-757 on figure(a)is caused by the gravity field of a luminous red galaxy which is distorting the light coming from a distant blue galaxy behind it (Credit: ESA/Hubble & NASA). Figure(b)shows the first planet detection using microlensing (Gaudiet al.,2008). Due to the duration of the microlensing event twelve telescopes were involved in order to have a continuous coverage of the event. The main peak is caused by the gravitational field of the host star while the small features numbered from 1 to 5 are caused by the two∼0.7 and∼0.27 MJupplanets.
its important mass, is generating astrong lens (Walshet al.,1979). Such a strong lens imaged byHubbleis represented on figure2.4a. Since the first detection of a microlens- ing event caused by a∼3 MJupplanet orbiting a 0.6M¯star (Bondet al.,2004), a total of 26 planets (3M⊕–10MJup, 0.2–8 AU) have been found with two multiple systems (Gaudi et al.,2008; Hanet al.,2013, figure2.4b).
The draw-back of this technique is that it is based on a single event detection. Once the planet is detected no direct follow-up or characterisation observation are possible.
When waiting long enough the host star moves away from the lens and becomes visible again for additional characterisation. Furthermore, the interpretation of a lensing event is not unique due to the complexity of the signal requiring approximations for the analy- sis (Han,2005), the significant number of parameters that can be adjusted to fit the data (Hanet al.,2013), and to the so-called close/wide degeneracy (Griest & Safizadeh,1998)
2.2 Direct Detection Techniques
Instead of indirectly searching for companions through effects they may have on the spectrum position or flux of the host star, direct detection methods aim at observing the flux directly emitted by the companion. The major difficulty of direct detection methods is the high contrast between the host and companion fluxes.
The straightforward method for direct detection is to subtract the flux from the host star in order to detect the flux emitted by the companion. The detection efficiency of this direct imaging method increases with orbital width. This is due to the fact that the light of the star forms a halo when being diffracted by the telescope (section3.1.1), and the increased separation decreases the effect of this stellar halo. One notable exception to this rule is the specific method which aims at detecting reflected light on the companion instead of thermal light where efficiency decreases with separation.
Due to the maximal orbital period coverage available currently as well as the de- crease of the induced signal, the radial velocity and transit techniques have a sharp cut-
2.2. DIRECT DETECTION TECHNIQUES 15
Figure 2.5– Mass to orbital separation diagram for detected sub-stellar companions, showing the limitation towards close orbits detection for radial velocity (squares) and transit (filled circles). The three outer planets of the solar system, Saturn, Uranus, and Neptune, are marked with triangles (Hinkley,2011).
off beyond 5–6 astronomical units (AU)(figure2.5). The direct detection techniques on the other hand are more sensitive at wider separation. These direct detection techniques are also applicable to a larger variety of stars from young active dwarfs to evolved mas- sive stars, as the intrinsic variability of the host is only a minor issue.
2.2.1 Direct Imaging
Even thoughtaking a pictureof a planet is probably the most evident method, it also turns out to be one of the most challenging techniques. The flux ratio between the star and a self-luminous giant planetary companion is at least 10−6in the favourable infrared domain. The separation of a companion with a semi-major axis of 5 AU at a rather small distance of 10 pc from Earth is 0.500and shrinks to 0.0500at 100pc, which is in both cases within the halo of the star. Using adaptive optics, coronography, and specific observ- ing methods it is nonetheless possible to detect planets. The direct detection of the three planets orbiting HR8799 (Maroiset al.,2008,2010) and a fourth one later on, along withβPictoris b (Lagrangeet al.,2009) launched the field of exoplanet direct imaging (Fig.2.6). A third planet discovery was announced at the same time orbiting Fomalhaut (Kalaset al.,2008), but its true nature remains enigmatic as it is only detected in the vis- ible but not in the infrared where a planet is expected to radiate most its light. A dust cloud is much more compatible with the observations than a planet (Marengo et al., 2009; Jansonet al.,2012b; Currieet al.,2013). These three detection and the nearly 50 following ones all made use of newly developed observing and data reduction methods, and a detailed description of them will be given in chapter3, as they are at the core of this work.
By comparing the directly imaged planets orbiting HR8799 with brown dwarfs of
16 CHAPTER 2. DETECTING SUB-STELLAR COMPANIONS
(a)HR8799 b, c, and d (b)βPic b
Figure 2.6– The two first confirmed planet direct detections. On(a)are the discovery images of HR8799 taken with Keck and GEMINI (Maroiset al.,2008), while(b)is a composite image showing the companion toβPic at two different observing epochs as well as the debris disk (Mouilletet al.,1997; Lagrangeet al.,2009,2010, composite picture courtesy of D.
Ehrenreich with the companion imaged with NACO/VLT, and the disk with ADONIS/ESO- 3.6m,)
analogousTeffwhich are expected to have similar atmospheric characteristics, it turned out that these planets were redder than their brown dwarf counterparts (Barmanet al., 2011, figure2.7). An explanation to this discrepancy given by Marleyet al.(2012) is that for a given temperature the low surface gravity of these young planets, induced dust clouds at higher altitudes than for brown dwarfs where these clouds would have pre- cipitated. This scenario has been further verified through observations of very young free-floating brown dwarfs with low surface gravity where the infra-red colours turned out to be similar to those of the directly imaged planets (Fahertyet al.,2013).
In some favourable cases a spectrum of the companion can be acquired. It has been tried on most directly detected companions with varying resolution depending on the contrast (Fig.2.8). The HR8799 system is a very favourable target for such observations due to the planet separations and luminosities. The spectrum acquisition is further fa- cilitated by the fact that the A-type host star has a nearly featureless spectrum so that the companion spectrum is not contaminated by features from the host spectrum. As a result, its companions have been intensively characterised yielding important results on their atmospheric chemistry and possible formation mechanism. The strong CO ab- sorption combined with the absence of methane observed inH andK band spectra are evidence for non-equilibrium chemistry in their atmospheres (Bowleret al.,2010;
Barmanet al.,2011). While the remarkable high resolution spectrum of HR8799c by Konopackyet al.(2013, figure2.9) made it possible to measure C and O abundances lower than expected. In addition to that, the derived C/O ratio they find is slightly higher than the one of the parent star which is all hinting towards a core-accretion scenario (Öberget al.,2011), but this planet is orbiting at 40 AU, far beyond the∼10 AU limit where this mechanism is supposed to stop working.
An important draw-back of the direct imaging technique is the important rate of false positives which is around 50% for current surveys targeting young stars (see8.2.2), mainly due to background M-dwarfs. A second epoch observation is therefore manda- tory to confirm that a detected candidate companion is bound. Precise astrometry as well as a long time-span between the two observations (usually more than a year) is
2.2. DIRECT DETECTION TECHNIQUES 17
Figure 2.7– The position of HR8799bcd and 2M1207b in near-IR colour–magnitude diagrams with respect to L and T field brown dwarfs (Leggettet al., 2002; Knappet al., 2004;
Casewellet al.,2007). The dotted lines show colour–magnitude tracks for chemical equi- librium models (R=1RJup, log(g)=4) of constant Teff (900 K to 1400 K in steps of 100 K) with cloud thickness varying from no clouds on the left to thick clouds on the right. The red and dusty planets are clearly off the L-T sequence. Figure from Barmanet al.(2011).
Vigan et al. (2008)
Close et al. (2005) Kasper et al. (2007) This work
Lafrenière et al. (2008)
Linear interpolation of stellar profile Symmetry of
stellar profile
Lyot coronagraphy+
spatial rescaling+
speckle fitting
Star Companion Speckle Slit Speckle/halo calibration zone
Tilted slit+
spatial rescaling+
speckle fitting Mohanty et al. (2007)
Moffat or Gaussian fitting of stellar profile
subtraction linear interpolation profile fitting fitting on rescaled speckles fitting on rescaled speckles
Spectral Deconvolution
Figure 2.8– The variety of high-contrast spectroscopy techniques, differing mainly in the way the stellar spectrum is subtracted from the companion one. The method used by the author of this illustration (Viganet al.,2012a) is represented on the far right.
needed in order to confirm that the companion is co-moving. And an even longer time- span is needed to observe the orbital motion of the companion, knowing that the cur- rently shortest orbital period detected is 20 years (βPic b).
2.2.2 Favoured Target Categories for Direct Imaging
In the past few years, many direct imaging surveys have been carried out using adap- tive optics on 8m-class telescopes with various target samples to optimise the scien- tific output. An intensively observed class of targets are young stars, typically younger than 1 Gyr and closer than 100 pc. Sub-stellar companions have no internal heat source which means that nearly all their thermal energy originates from their formation, so that the thermal energy lost through radiation is not replaced. This evolution is represented on figure2.10which shows the luminosity as a function of age for the Baraffeet al.(2003)
18 CHAPTER 2. DETECTING SUB-STELLAR COMPANIONS
Figure 2.9– K-band spectrum of the 3-7 MJup companion HR 8799c (black) acquired with OSIRIS/Keck II by Konopackyet al.(2013). This observations resulted in the detection of carbon monoxide and water vapour absorption lines, C/O ratio and abundances hint- ing towards a formation through core-accretion.
and Marleyet al.(2007) evolutionary models. Their young age make it possible to detect low mass companions which are still bright as they didn’t have time to cool down. These prime targets have been observed in numerous surveys, such as Masciadriet al.(2005), B. A. Billeret al.(2007), Kasperet al.(2007), Lafrenièreet al.(2007), Apaiet al.(2008), Chauvinet al.(2010), Heinzeet al.(2010), Jansonet al.(2011), Delormeet al.(2012), Viganet al.(2012b), B. A. Billeret al.(2013b), Jansonet al.(2013), Nielsenet al.(2013), Rameauet al.(2013a), Wahhajet al.(2013), Yamamotoet al.(2013), Chauvinet al.(2015), and Desideraet al.(2015). Moving group members are an important sub-category of young stars (10–300 Myr) with the advantage of having a well constrained age (Kraus et al.,2008,2011; Evanset al.,2012). Some surveys specifically target debris disk host- ing stars achieving a slightly higher detection rate (Apaiet al.,2008; Jansonet al.,2013;
Wahhajet al.,2013), and the first two stars around which exoplanets were detected also hosted debris disks (Maroiset al.,2008; Lagrangeet al.,2009). A survey of 60 G stars with theHerschelsatellite revealed that stars hosting planets lower than saturn mass tend to have a higher frequency of detectable debris disks (Wyattet al.,2012). Aiming at less luminous targets such M, white, and brown dwarfs makes it possible to reach smaller planets thanks to the reduced contrast between the host and the companion (Burleigh et al.,2002; Bowleret al.,2012; Delormeet al.,2012), but extrapolation from radial ve- locity surveys have shown that planet occurrence increases with stellar mass (Johnson et al.,2010; Bonfilset al.,2013).
Another approach is to aim at stars where something indicates the presence of com- panions within the detection range. Results from radial velocity surveys show that stars are more likely to host multiple planets than single planets (Udry & Santos,2007) which motivated the yet unsuccessful search by direct imaging of wide companions around
2.3. COMBINING THE TECHNIQUES FOR ADDITIONAL CHARACTERISATION 19
(a)Lyon models (b)Tucson models
Figure 2.10– Luminosity as a function of time for jupiters of various masses. Young objects are more favourable for detection due to their thermal brightness. Figure(a)is based on the DUSTY and COND models (Baraffeet al.,2003), while figure(b)shows the evolutionary tracks from the Tucson models for hot start and core-accretion, dashed and solid lines respectively (Marleyet al.,2007).
planet hosting stars (Jenkinset al.,2010). Furthermore a massive companion on a wide orbit can perturb the orbit of the inner objects though Kozai cycles (Kozai,1962; Fab- rycky & Tremaine,2007; Nagasawaet al.,2008; Katzet al.,2011) which is one mechanism to form hot jupiters and misaligned companions. Direct imaging campaigns targeting such objects were carried out without any companion detection yet (Naritaet al.,2010).
Using radial velocity it is also possible to detect signs of a wide orbit companion through the long-term drift in the data. By directly imaging such targets it is then possi- ble to directly detect the companion, or give an upper mass limit in case of non-detection.
This approach is the one followed in this work (chapter4–6), as well as by the TRENDS survey (Creppet al.,2012).
2.3 Combining the Techniques for Additional Characterisation
The detection methods presented above are not only complementary in order to cover most of the mass–separation parameter space, they also give access to different infor- mation on the companions detected. Radial velocity detections provide information on lower mass limits of the companion (due to the unknown inclinationMsini), as well as eccentricity and semi-major axis. Detecting companions through transit provides the radius, and in some favourable cases a glance into atmospheric day to night side circu- lation (Charbonneauet al.,2005). Transmission spectroscopy observations of transiting planets are aiming at detecting atmospheric composition of a few ideal targets (Char- bonneauet al.,2002). Combining radial velocity and transit detection solves the mass degeneracy as the inclinationiis known, and gives access to the bulk densityρsince the mass and radius are known. The radial velocity pseudo variation observed during a tran- sit gives the alignment of the companion orbit with respect to the stellar spin (McLaugh- lin,1924; Rossiter,1924; Quelozet al.,2000; Winnet al.,2005). Using astrometric data from the Hipparcos satellite Sahlmannet al.(2011,2013) were able to determine the true mass of known radial velocity companions.
Direct imaging detection provides projected separation, photometry, and sometimes even emission spectroscopy of the companion (Bonnefoyet al.,2010; Konopackyet al., 2013), but the mass has to be extrapolated from the photometry. When combined with astrometry or radial velocity one gets the mass of directly imaged companions.
C H A P T E R
3
A T ECHNIQUE TO D IRECTLY I MAGE C OMPANIONS ON W IDE O RBITS
Chapter contents
3.1 Direct Imaging of Companions . . . 20 3.2 GRAPHIC: The Geneva Reduction and Analysis Pipeline for High-contrast
Imaging of planetary Companions . . . 27 3.3 Publications on GRAPHIC . . . 29 3.4 Further improvements . . . 40 In this chapter I will first present the direct imaging technique, followed by a de- scription of the data reduction pipeline, used for our observing program presented in chapter4. The first published result on companion search and disk imaging based on this pipeline is included in chapter9.
3.1 Direct Imaging of Companions
Most direct imaging observations are based on the combinations of three instrumental techniques. Adaptive optics to compensate for the atmospheric turbulence, coronogra- phy for a higher contrast ratio, and differential imaging to subtract the diffracted light of the host star from that of the companion. Ground-based observations suffer from the wavefront distortion induced when the light travels through the atmosphere. The use of adaptive optics is thus mandatory for ground-based direct imaging to achieve a high enough spatial resolution and to increase the signal of any possible companion by sharpening its point-spread-function (section3.1.1). Coronography can be used to reach higher contrast ratios between the star and the companion flux by masking the light from the parent star on the detector (section3.1.2). Finally, residual diffracted light from various aberrations can be suppressed by means of differential imaging (section 3.1.3and3.1.4).
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