Introduction: planet detection and studies in the historical context

Beyond Earth, only five planets have been known since historical times. These are the three “terrestrial planets,” Mercury, Venus and Mars, and the two “gas giants,” Jupiter and Saturn. The systematic studies of the skies starting in the seventeenth century marked the beginning of a new era. The ice giant Uranus, the next outer planet beyond Saturn, was first cataloged as 34 Tau by John Flamsteed in 1690, but not recognized as a “wanderer” in the skies. In 1781 William Herschel was the first to spatially resolve the disk of Uranus, initially classifying it as a comet. Uranus also played an important role in the discovery of Neptune. Neptune, the eighth planet in the Solar System, and the outermost of the ice giants, was first observed in 1612/1613 by Galileo Galilei. At that time, Neptune was in close conjunction with Jupiter. Galileo recognized the moving source in the vicinity of Jupiter (see Fig. 16.1), but decided not to get distracted from his studies of the orbital motions of Jupiter's four largest moons, which we now refer to as “Galilean Moons.” Another 230 years passed till precise measurements of Uranus' orbit indicated the presence of an additional outer planet. Based on predictions by Urbain Leverrier (which were in close agreement with independent calculations by John Couch Adams) in 1846, Johann Gottfried Galle finally identified and resolved Neptune.

Introduction

At the time of this workshop, there are now more than 150 detected extrasolar planets discovered and 13 confirmed multiple planet systems with more candidate planets and systems being evaluated (G. Marcy in Chapter 11 of this book). There is a growing number of ground-based observations of young circumstellar and protoplanetary disks as well as volumes of data from Spitzer and Cassini. There are two scenarios for gas-giant planet formation that are sufficiently sophisticated to provide results and predictions. Clearly, the ingredients are present for planetary scientists to develop a comprehensive or, at least, a cohesive model for the formation of the gas-giant planets in the Solar System and in extrasolar systems and, to initiate resolution of the age old question: how do planets form?

The subject of this chapter is the core accretion–gas capture model, the more generally accepted scenario for gas-giant planet formation (see also Chapter 8 by Thommes and Duncan for the core accretion itself). Please note that this model has had many names over the decades of its development: planetesimal hypothesis, nucleated instability model, and core instability model. However, it seems best to call it by the more descriptive label: the Core Accretion–Gas Capture model, the CAGC or, the short version, the core accretion model.

Introduction

The two mechanisms that have been advanced for explaining the formation of giant planets are core accretion (“bottom up”) and disk instability (“top down”). Core accretion, the conventional mechanism, relies on the collisional accumulation of planetesimals to assemble ∼ 10 M⊕ solid cores, which then accrete massive gaseous envelopes from the disk gas (Mizuno, 1980; Lissauer, 1987; Pollack et al., 1996; Kornet et al., 2002; Inaba et al., 2003). In this scenario, the ice-giant planets probably did not form in situ (Levison and Stewart, 2001), but rather formed between Jupiter and Saturn and then were scattered outward to their present orbits (Thommes et al., 1999, 2002).

The alternative to core accretion is disk instability, where gas-giant protoplanets form rapidly through a gravitational instability of the gaseous portion of the disk (Cameron, 1978; Boss, 1997, 1998b, 2000, 2002a,b, 2003, 2004; Gammie, 2001; Mayer et al., 2002, 2004; Nelson, 2000; Nelson et al., 1998, 2000; Pickett et al., 1998, 2000a,b, 2003; Rice and Armitage, 2003, Rice et al., 2003a) and then more slowly contract to planetary densities. Solid cores form simultaneously with protoplanet formation by the coagulation and sedimentation of dust grains in the clumps of disk gas and dust.

Recent improvements in radial velocity precision

Since the first discovery of an extrasolar planet around a Solar-type star ten years ago (Mayor and Queloz, 1995), research in this field has been very productive and has led to the detection of more than 140 exoplanets. The vast majority of these discoveries has been made with the radial-velocity (RV) technique, i.e. the precise measurement of the RV wobble that a planet induces in its parent star due to its orbital movement. A major effort to improve the accuracy of the RV measurements has been undertaken by several groups, since this is absolutely necessary to detect the RV signatures of giant planets, in the range 1–100 ms−1. Two main techniques were developed: one using a ThAr calibration simultaneously with each observation (Baranne et al., 1996) to track instrumental drifts, and one using an iodine absorption cell, superimposing a reference spectrum on the stellar spectrum (Butler et al., 1996). Both techniques have been able to deliver RV precision at the level of ∼3 ms−1, opening the way to the discovery of many planetary systems.

Over the past decade, the exoplanet group at Geneva Observatory has been operating two high-resolution spectrographs able to achieve high RV precision, namely the ELODIE instrument mounted on the 1.93 m telescope at Observatoire de Haute-Provence (France), and the CORALIE instrument installed on the Swiss 1.2 m telescope at La Silla Observatory (Chile). Both ELODIE and CORALIE are high-resolution (R = 50 000), fiber-fed echelle spectrographs.

Introduction

Planets form within circumstellar disks composed of a mixture of gas and dust grains. These disks result from the gravitational collapse of rotating molecular cloud cores. They are initially rather massive and consist of about 0.3 M*, where M* is the mass of the central star (e.g. Yorke et al., 1995). In contrast, the minimum mass required to build the planets of our Solar System is only about 0.01 Solar masses (M⊙). Evidently, there are processes that redistribute the mass, transform the dust to larger particles, and disperse much of the gas and dust.

The processes which are responsible for the dispersal of the gas influence the formation of planets. For example, the timescale for gas dispersal as a function of the disk radius affects the composition of the resulting planetary system. As long as the dust particles are small enough to be tightly coupled to the gas, they follow the gas flow. If the gas is dispersed before the dust particles have had a chance to grow, all the dust will be lost and planetesimals and planets cannot form. Even if there is time for particles to coagulate and build sufficiently large rocky cores that can accrete gas (Pollack et al., 1996; Hubickyj et al., 2004), the formation of gas-giant planets like Jupiter and Saturn will be suppressed if the gas is dispersed before the accretion can occur.

Introduction: geoscience meets astronomy

Much of our knowledge about the formation of planets in the Solar System and in particular concepts and ideas about the origin of the Earth are derived from studies of extraterrestrial matter. Meteorites (Sears, 2004; Lauretta et al., 2006; Krot et al., 2006) were available for laboratory investigations long before space probes were sent out for in situ investigations of planetary surfaces, or Moon rocks were brought back to Earth. Meteorite studies provided such important parameters as the age of the Earth and the time of formation of the first solids in the Solar System (Chen and Wasserburg, 1981; Allegre et al., 1995; Amelin et al., 2002), as well as the average abundances of the elements in the Solar System (Anders and Grevesse, 1989; Palme and Jones, 2004). Traditionally, the study of rocky material requires techniques that fundamentally differ from astronomical techniques. While electromagnetic radiation from stars is analyzed by spectroscopy, the solid samples of aggregated cosmic dust and rocky matter from planetary surfaces require the use of laboratory instruments that allow the determination of their chemical and isotopic composition. Planetary surface materials are present in polymineralic assemblages. Formation conditions and thermal stability of individual minerals provide important boundary conditions for the genesis and history of the analyzed materials. Such studies require a thorough mineralogical background. The abundances and properties of the rock-forming elements, the focus of geo- and cosmochemical research, are, however, not necessarily of major concern to astrophysicists.

Introduction

Given that brown dwarfs are usually much more massive than planets (see Section 15.6), it is somewhat surprising that the first incontrovertible discovery of a brown dwarf (Nakajima et al., 1995) and the discovery of the first extrasolar planet (Mayor and Queloz, 1995) were announced simultaneously in 1995. Over the past decade, the rapid progress made in both fields has been extraordinary. There are now more than 150 extrasolar planets known, including more than a dozen multiple-planet systems. The first brown dwarf, Gliese 229B, was found in orbit around an M-dwarf, but in the same year other candidates, later confirmed to be free-floating brown dwarfs, were announced (e.g. Teide 1 by Rebolo et al., 1995), along with PPl 15 which was later discovered to be a binary brown dwarf (Basri and Martin, 1999). Observations now suggest that brown dwarfs are as common as stars, although stars dominate in terms of mass (e.g. Reid et al., 1999).

Since the rest of this book is devoted to the topic of planets, in this chapter I will review the properties and potential formation mechanisms of brown dwarfs, comparing and contrasting them with planets, but referring the reader to the other chapters of the book for detailed information on planets.

Masses

The most fundamental parameter of a brown dwarf is its mass.

Introduction

According to the core-accretion model for planet formation, the building blocks of planets are formed by the coagulation of dust grains, growing from the initial sub-micron sizes inherited from the interstellar medium to the 100 kilometer sizes of full-grown planetesimals. This is a growth process over 12 orders of magnitude in size and 36 orders of magnitude in mass. The physics of dust is of crucial importance for the study of planet formation. It also plays a major role in the structure and evolution of protoplanetary disks, since the dust carries most of the opacity of the dust–gas mixture in these disks and provides the surface for chemical reactions. Moreover, infrared and (sub)millimeter observations of dust continuum emission from these disks can be used as a powerful probe for the disk structure and mineralogical composition. A deep understanding of the physics of dust and the coagulation of grains is therefore of paramount importance for the study of the formation of planets and the circumstellar disks in which they are formed.

The study of grain coagulation and the formation of planetesimals has a long history. At the start of the twentieth century an equation for coagulation of colloidal particles was formulated by Smoluchowski (1916), though not related to astrophysical applications. The continuous form of that equation was later used to study the size distribution of fog particles in the Earth's atmosphere (Shumann, 1940).

Introduction

Rather few facts can be considered as acceptable to all who are working in the field of planet and planetesimal formation. Starting there, we will explore the possible pathways as suggested by experiments. It is certainly undisputed that the regular mode of planet formation is connected to protoplanetary disks. These disks consist mostly of gas, which makes up about 99% of their mass. The remaining 1% resides in the form of dust and – depending on the temperature – in the form of ice. As terrestrial planets are mostly built from heavier elements it is natural to assume that they are somehow assembled from the dust component in the disk.

Whatever model is placed between the dust and the planets, collisions between the solid bodies are unavoidable. In fact a large part of the process of planet formation can be based on collisions which can and (at least partly) will lead to the formation of larger bodies.

In the following sections we will review experiments that have studied these collisions and eventually put these results in a rough sketch of planetesimal formation. It is sometimes argued that collisions of large bodies might be too energetic to lead to the formation of a still larger body (Youdin and Shu, 2002). As described in this chapter it is true that collisions can lead to erosion rather than growth. However, we will show that this is not necessarily so for all collisions.

Introduction

The observation of a transit in our own Solar System is a long-lasting experience. Historical events related to transits can be traced back to Ptolemy who mentioned in his “Almagest” that the lack of detections of transits was not in contradiction with Mercury and Venus being closer to the Earth than the Sun (in the geocentric system) simply because they could be either too small to be detected or their orbital plane could be slightly tilted to the Solar one (Gerbaldi, personal communication). In 1607 Johannes Kepler thought he had directly observed a predicted Mercury transit but in fact only followed sun spots. He did, however, predict the next transits of Venus and Mercury to take place in 1631 following the extremely accurate observations of the planets by Tycho Brahe. The first transit to be observed was the Mercury transit in 1631 with the best observations leading Pierre Gassendi to evaluate its diameter to be less than 20 arcsec, much smaller than ever thought before. All the following transit observations led to new ephemerides and estimates of the size of the Solar System, but not as accurate as expected because of the difficulty of locating in time the entrance and exit of the planetary disk over the Solar one. First pictures of the Venus transit were made as early as in 1874 (Fig. 9.1). The transit of Mercury was also observed with the Solar and Heliospheric Observatory (SOHO) spacecraft from the L5 Lagrange point of the Earth.

Introduction

The observed characteristics of molecular clouds from which stars form can be reproduced by simulations of magnetohydrodynamic (MHD) turbulence, indicating the vital role played by magnetic fields in the processes of star formation. The fields support dense cloud cores against collapse, but they cannot do so indefinitely, because only charged particles couple to the field lines while neutral atoms and molecules can freely slip through. Through this process, called ambipolar diffusion, the cores slowly contract. The recombination rate in denser gas increases, causing the ionisation degree of the core to decrease. According to available observational data, once the core has contracted to ∼0.03 pc it decouples from the magnetic field and enters the dynamic collapse phase. During the collapse the angular momentum is locked into the core and remains unchanged (Hogerheijde, 2004).

Protostellar collapse and formation of disks

The typical specific angular momentum of a core on the verge of dynamic collapse, jc, amounts to ∼1021 cm2 s−1, and is many orders of magnitude larger than the typical specific angular momentum of a star (Hogerheijde, 2004). The inevitable conclusion is that the protostellar object resulting from the collapse must be surrounded by a large, rotationally supported disk (hereafter, protoplanetary disk) in which the original angular momentum of the core is stored. The outer radius of the disk, rd, may be roughly estimated based on Kepler's law.

This conference truly reflects a microcosm of an explosive revolution in the quest to understand the origin of planet and star formation. The diverse nature of this wide-open field necessitates a multi-facet attack on all relevant issues. In this pursuit, it is particularly important to find the missing links between the many seemingly independent observations as circumstantial clues around a global picture. The development of a comprehensive coherent interpretation requires an integrated approach to identify the dominant physical processes which determine the physical characteristics of planets and the dynamic architecture of planetary systems.

On the basic concept of planetary origin, there is very little difference between modern theories and the original Laplacian hypothesis. The coplanar geometry of all the major planets' orbits hardly needs any extrapolation for theorists to postulate the scenario that the planets formed long ago in a rotational flattened disk which is commonly referred to as the Solar Nebula. Today, we have direct images and multi-wavelength spectra of protostellar disks within which planet formation is thought to be an ongoing process. Perhaps the biggest advancement in the past decade is the discovery of over 100 planets around nearby stars other than the Sun. For the first time in this scientific endeavor, the Solar System reduces its unique importance to a single entry in the rapidly growing database of planetary-system census.

The feedback effects of massive stars on their galactic and intergalactic environments can dominate evolutionary processes in galaxies and affect cosmic structure in the Universe. Only the Local Group offers the spatial resolution to quantitatively study feedback processes on a variety of scales. Lyman continuum radiation from hot, luminous stars ionizes H II regions and is believed to dominate production of the warm component of the interstellar medium (ISM). Some of this radiation apparently escapes from galaxies into the intergalactic environment. Supernovae and strong stellar winds generate shell structures such as supernova remnants, stellar wind bubbles, and superbubbles around OB associations. Hot (106 K) gas is generated within these shells, and is believed to be the origin of the hot component of the ISM. Superbubble activity thus is likely to dominate the ISM structure, kinematics, and phase balance in starforming galaxies. Galactic superwinds in starburst galaxies enable the escape of mass, ionizing radiation, and heavy elements. Although many important issues remain to be resolved, there is little doubt that feedback processes plays a fundamental role in energy cycles on scales ranging from individual stars to cosmic structure. This contribution reviews studies of radiative and mechanical feedback in the Local Group.

Introduction

The Local Group is especially suited as a laboratory for studying the effects of the massive star population on the galactic environment. There are three types of massive star feedback:

1. (a) Radiative feedback, i.e., ionizing emission, which results in photoionized nebulae and diffuse, warm (104 K) ionized gas;

2. (b) Mechanical feedback, predominantly from supernovae (SNe), resulting in supernova remnants (SNRs), superbubbles, and galactic superwinds; and

3. […]

The Galaxy's extended halo contains numerous satellites which are in the process of being disrupted. This paper discusses the stages of satellite accretion onto the Galaxy with a focus on the Magellanic Clouds and Sagittarius dwarf galaxy. In particular, a possible gaseous component to the stellar stream of the Sgr dwarf is presented that has a total neutral hydrogen mass between 4–10×106 M⊙ at the distance to the stellar debris in this direction (36 kpc). This gaseous stream was most likely stripped from the main body of the dwarf 0.2–0.3 Gyr ago during its current orbit after a passage through a diffuse edge of the Galactic disk with a density > 10−4 cm−3. This gas represents the dwarf's last source of star formation fuel and explains how the galaxy was forming stars 0.5–2 Gyr ago. This is consistent with the star formation history and H I content of the other Local Group dwarf galaxies.

Introduction

Our Galaxy has built itself up by accreting satellite galaxies. This process if evident today through the satellites currently found in the extended Galactic halo. There are nine satellite galaxies within 150 kpc interacting with our Galaxy at various levels. These are in order of distance (Grebel, Gallagher, & Harbeck 2003): the Fornax dSph (138 kpc), the Carina dSph (94 kpc), the Sculptor dSph (88 kpc), the Sextans dSph (86 kpc), the Draco dSph (79 kpc), the Ursa Minor dSph (69 kpc), the Small Magellanic Cloud (63 kpc), the Large Magellanic Cloud (50 kpc), and the closest example of a recognizable accreting satellite is the Sagittarius Dwarf (28 kpc; hereafter Sgr dwarf).

There is increasing observational evidence that hot, highly ionized interstellar and intergalactic gas plays a significant role in the evolution of galaxies in the local universe. The primary spectral diagnostics of the warm-hot interstellar/intergalactic medium are ultraviolet and X-ray absorption lines of O VI and O VII. In this paper, I summarize some of the recent highlights of spectroscopic studies of hot gas in the Local Group and low-redshift universe. These highlights include investigations of the baryonic content of low-z Ovi absorbers, evidence for a hot Galactic corona or Local Group medium, and the discovery of a highly ionized high velocity cloud system around the Milky Way.

Introduction

We live in a wonderful age of discovery and exploration of the universe. As we peer farther and farther back in time, it is becoming ever more important to make sure that we observe the local universe as well as possible. Observations of galactic systems and the intergalactic medium (IGM) in the low-redshift universe are required to study the universe as it has evolved over the last ∼5 billion years. They are essential for the interpretation of higher redshift systems, and they form a framework for studies of such key topics as galactic evolution, “missing mass,” and the distribution of dark matter. Studies of hot gas and its relationship to galaxies are shedding new light on these and other astronomical topics of interest today. In this review, I summarize some basic information about the elemental species and types of observations that can be used to study hot gas.

The primordial abundances of deuterium, helium-3, helium-4, and lithium-7 probe the baryon density of the Universe only a few minutes after the Big Bang. Of these relics from the early Universe, deuterium is the baryometer of choice. After reviewing the current observational status of the relic abundances (a moving target!), the baryon density determined by big bang nucleosynthesis (BBN) is derived. The temperature fluctuation spectrum of the cosmic background radiation (CBR), established several hundred thousand years later, probes the baryon density at a completely different epoch in the evolution of the Universe. The excellent agreement between the BBN- and CBR-determined baryon densities provides impressive confirmation of the standard model of cosmology, permitting the study of extensions of the standard model. In combination with the BBN- and/or CBR-determined baryon density, the relic abundance of 4He provides an excellent chronometer, constraining those extensions of the standard model which lead to a nonstandard early-Universe expansion rate.

Introduction

As the hot, dense, early Universe rushed to expand and cool, it briefly passed through the epoch of big bang nucleosynthesis (BBN), leaving behind as relics the first complex nuclei: deuterium, helium-3, helium-4, and lithium-7. The abundances of these relic nuclides were determined by the competition between the relative densities of nucleons (baryons) and photons and, by the universal expansion rate. In particular, while deuterium is an excellent baryometer, He provides an accurate chronometer.