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4 Extrasolar planets: detection, properties,
projects
Abstract:
Detection methods: Many complementary detection methods exist. The easiest
methods are indirect and exploit the influence of planets on their host star, e.g. the most
successful radial velocity method, astrometry, or the transit method. These methods
allow us to study basic parameters of planets (orbital parameters, mass, radius, density)
More challenging methods aim at directly observing the light from planets, e.g. the
polarization and direct imaging method. In future, they will allow us to study in depth
the physics of planets (composition, surface structure, signatures of life…).
Properties of extrasolar planets: Currently over 400 exoplanets have been detected,
mainly Jupiter-like giants due to observational biases. Extrasolar planets (detected so
far) are more common around metal-rich host stars, as expected from the standard
planet formation theory. Many Jupiters were found at very small orbits with periods of
3–4 days (hot Jupiters), which has led to the new concept of planet migration (to allow
Jupiters to form outside the ice-line and then migrate inward to very small orbits).
Projects: Many projects are planned in the near future, both ground-based and in space.
Within 20 years it will be possible to detect ten thousands of planets including Earthsized terrestrial planets (even at 1 AU orbital radius as Earth) and search for signatures
of life.
4.1 Detection methods
Radial velocity
A planet orbiting a star causes the star to rotate around their common center of mass.
This can be detected as a Doppler shift in the star's spectral lines. The star's spectrum is
measured over time and a periodic shift of
spectral lines then indicates an orbiting body. The
light originating from a star moving towards Earth
will be Doppler shifted to bluer (shorter)
wavelengths, while a star receding from Earth will
emit light shifted to redder (longer) wavelengths.
The effect is very small. For instance Jupiter
induces a 12 m/s velocity change of the Sun
whereas the effect by Saturn is only 2.7 m/s. The
most sensitive instrument at the moment can
Figure 1: Radial velocity method.
detect Doppler shifts of a bit less than 1 m/s.
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This method is most sensitive to heavy planets with a plane of rotation parallel to our
line of sight and small orbital periods. A major drawback is that the angle of inclination
(angle between the direction perpendicular to the plane of rotation and our line of sight)
is usually not known, so that there is a high uncertainty in the derived mass of the
planet. The determined value of the mass is actually Msin(i), where i is the inclination.
Figure 2: Orbital motion of 51 Peg. The solid line represents
the computer model fitted to the data. The orbital motion is
due to a planet of about half Jupiter’s mass with an orbital
period of 4.2 days. This was the first detection of an
extrasolar planet around a Sun-like star (Mayor & Queloz
1995).
Almost all extrasolar planets that have been found so far have been detected with the
radial velocity method.
Let us briefly look at how the orbital radius, the planetary mass, or rather Msin(i), and
the eccentricity can be determined with this method.
The orbital radius of the planet ap is obtained with Kepler’s third law
 a3 
P 2  4 2  p  ,
 GM 
*

where P is the orbital period, G = 6.67  1011 Nm2kg2, and M* is the mass of the
central star. The orbital period is easily measured due to the periodic variations of the
radial velocity (cf. Fig. 1). The stellar mass can be obtained from stellar evolution
models, or in the case of main-sequence stars from the luminosity-mass relation L  M4
for M > 0.4M, and L  M2.8 for M < 0.4M).
The planetary mass is found from the fact that the two bodies (star and planet) revolve
about their common center of mass. The distances, or radii, to this common center of
mass are related to the masses by
M *a*  M p ap .
The stellar mass M* and the orbital radius ap of the star are already known. But we also
require a* to determine the mass of the planet. In a circular orbit, the orbital speed v* is
constant and the period P is defined as the time taken to complete one orbit. This means
that the distance traveled in one orbit must be equal to the circumference of the circle,
2a*, so that
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v* P  2 a* .
However, we can usually not easily determine v* because the inclination i of the orbit is
not known, or only difficult to measure precisely. The inclination is defined as the angle
between the line-of-sight and the direction perpendicular to the orbital plane (i.e. i = 0°
if seen face-on; i = 90° if seen edge-on). Therefore, from radial velocity observations
(e.g. Fig. 1) the maximum radial velocity indicates only v*sin(i), the projection to the
line-of-sight, and a lower limit on the actual speed of the star. Nonetheless, this allows
us to set a lower limit to the star’s orbital radius
a* sin  i  
v* sin  i  P
.
2
As a consequence we find with above equations the lower limit on the planet’s mass
M p sin  i  
a* sin  i 
M* .
ap
Above we have assumed a circular orbit. In principle we could also assume an elliptic
orbit and follow the same basic idea. The eccentricity of the planetary orbit is obtained
from the shape of the radial velocity curve (Fig. 2), which then obviously does not
follow anymore a sin-law.
Transits
Photometry measures the intensity of the light of a single star accurately. If a planet
transits through the line of sight, this shows up as a tiny dip in the star’s intensity,
typically in the order of one percent (much smaller differences are difficult to detect
with today’s instruments). It is best compared to a partial eclipse of our own sun by
either Venus or Mercury. So, if this intensity dip occurs periodically, this could very
well indicate a planet.
Figure 3: Illustration of the light curve during a primary transit of a planet.
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Advantages of this method are the possibility to detect small planets and its
independence from the extrasolar planet’s orbital radius. The main advantage of the
transit method is that the size of the planet can be determined from the lightcurve. When
combined with the radial velocity method (which determines the planet's mass) one can
determine the density of the planet, and hence learn something about the planet's
physical structure.
The radius of the planet is obtained from the drop in flux (during the transit) and the
Radius of the star:
 Rplanet 
f  

 Rstar 
2
.
The transit method also makes it possible to study the atmosphere of the transiting
planet. When the planet transits the star, light from the star passes through the upper
atmosphere of the planet. By studying the high-resolution stellar spectrum carefully, one
can detect elements present in the planet's atmosphere. A planetary atmosphere (and
planet for that matter) could also be detected by measuring the polarisation of the
starlight as it passed through or is reflected off of the planet's atmosphere. Additionally,
the secondary eclipse (when the planet is blocked by its star) allows direct measurement
of the planet's radiation. If the star's photometric intensity during the secondary eclipse
is subtracted from its intensity before or after, only the signal caused by the planet
remains. It is then possible to measure the planet's temperature and even to detect
possible signs of cloud formations on it.
This method has two major disadvantages. First of all, planetary transits are only
observable for planets whose orbits happen to be perfectly aligned from astronomers'
vantage point. About 10% of planets with small orbits have such alignment, and the
fraction is far smaller for planets with larger orbits. However, because transit surveys
can scan large areas of the sky at once, the probability of finding extrasolar planets
could potentially exceed that of the radial-velocity method.
Figure 4: Physical quantities that can be studied by observing transits. Detections of planets by transits
often have to be confirmed by radial velocity measurements, because a transit could be mistaken for star
spots due to stellar activity. But once confirmed transits in principle allow us to obtain much more
complementary information on the planet.
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Secondly, the method suffers from a high rate of false detections. A transit detection
requires additional confirmation, typically from the radial-velocity method.
Microlensing
Microlensing makes use of Einstein’s notion of the curvature of space. In the early 20th
century, Einstein discovered that gravity causes space to curve, similar to the bending of
a bridge under its own weight. Einstein postulated that light that travels through space
according to this curved path. An immediate consequence is that any heavy object in
space could function as a gravitational lense. Even objects that are as heavy as stars can
function as a microlense. When one star passes in front of the other, an increase in
intensity due to the microlensing effect can be observed. If such a star harbors a planet,
this planet will increase the microlensing effect of a star. Such observations are possible
by observing many stars simultaneously – like the OGLE project does – and measuring
the flux rapidly.
Figure 5: Principle of microlensing and observational signature of planet. Left panel: A foreground star
focuses the light of a background star if the two stars are perfectly aligned as seen from Earth. The
resulting “Einstein ring” cannot be resolved in the case of microlensing, but it leads to an temporal
increase of the light curve of the background star when the foreground star passes through. Right panel: If
the foreground star harbors a planet, the planet itself increases the lensing effect resulting in a narrow
peak in the light curve.
Some advantages of microlensing are that it is a sensitive method so that you do not
need to wait a long period of time. Furthermore, the host star can be or even preferably
is to be faint; lastly, it is currently a promising method for detecting terrestrial
exoplanets using ground-based telescopes.
Figure 6: Observed microlensing event with
the characteristic peak caused by a planet.
The planet has been identified as a 5.5 Earth
mass planet. (Beaulieu et al. 2006)
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A notable disadvantage of the method is that the lensing cannot be repeated because the
chance alignment never occurs again. Also, the detected planets will tend to be several
kiloparsecs away, so follow-up observations with other methods are usually impossible.
Polarization
The polarization method searches for the light of the central star that is scattered on the
surface of the planet. The scattered light is linearly polarized, perpendicular to the
scattering plane. Therefore, the direction of the linear polarization rotates as the planet
orbits around the central star. The observed polarization signature thus exhibits the
orbital period of the planet, and, using Kepler’s third law, the orbital radius. If the orbit
is inclined the degree of polarization would vary, which allows us to determine the
inclination of the orbit. The eccentricity of the orbit can also be obtained.
Figure 7: Direction and degree of polarization expected of the light scattered by a planet orbiting its
central star at two different inclinations.
The light coming directly from the star could in general also be polarized, due to surface
magnetic fields. This results in possible false detections similar to the transit method so
that a confirmation by the radial velocity method would be ideal. The stellar
polarization signal can be distinguished from a planet, because the stellar signal would
not be perfectly periodic over long time scales. If the rotational period of the star differs
from the orbital period of the planet, then polarization from the planet would also have
different periodicity.
The polarization method has tremendous potential and advantages. In particular, it is not
an indirect method (as e.g. the radial velocity method that observes only the light of the
central star) since it allows us to observe photons coming from the planet itself.
Therefore, one could find the composition of the atmosphere, the albedo, and variations
of the albedo over the planetary surface (from the time variation of the polarization
signal), and thus the surface structure. This is a property shared with the direct imaging
method. In the polarization method it is however not necessary to actually spatially
resolve the planet (as in direct imaging), a big advantage for this method.
Linear polarization from an extrasolar planet might have been detected in the case of
HD189733 b, but the result is still controversial.
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Astrometry
Astrometry uses the position of a star to detect planets and uses the idea that a planet
around a star will cause a small "wobble". This movement is of course also used in the
radial velocity method. In this case, however, astronomers are searching directly for the
tiny displacements of the stars on the sky. Basically, the position of the central star has
to be determined to high precision, and followed over a long time, so that the wobbling
of the star becomes, which makes it possible to determine the planet’s period and mass.
The best method to do accurate astrometry is
to take a very far away (so that it won't move)
reference star at a small angle from the star,
which is to be observed. In order to apply this
method, more accurate instruments, both on
earth and in space, are currently being
developed.
Figure 8: Astrometric displacement of the Sun (almost
1 milli-arcsecond) due to the massive planets in our
solar system as it would be observed from 10 parsecs,
or about 33 light-years.
Direct imaging
Direct imaging is a technique that images planets directly. It uses the starlight reflected
by the planet to make it visible on a CCD. Ideally, the method uses a coronagraph (or,
more practically, a nulling interferometer) to block the light from the star in order to see
whether other light sources can be detected close to the star; if such a faint light source
is found, other methods are used to check whether it really is a planet.
Direct imaging is very difficult with current instruments. But it would of course allow
us to study a wealth of physical properties far beyond just the physical parameters of its
orbit: abundances, composition, temperature, surface structure, rotation period…
Figure 9: Possible direct detection of an extrasolar
planet. The white (larger) object is a brown dwarf
(2M1207 a). The faint red object 2M1207 b is 100
times fainter, intrinsically, than the bright brown dwarf
2M1207a (white in the picture)  a characteristic well
explained by a planet roughly five times the mass of
Jupiter. There is however ongoing controversy whether
2M1207 b is really a planet or also a low-mass brown
dwarf.
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Pulsar timing
Pulsars (the small, ultra dense remnant of a star that has exploded as a supernova) emit
radio waves extremely regularly as they rotate. Slight anomalies in the timing of its
observed radio pulses can be used to track changes in the pulsar's motion caused by the
presence of planets. Like an ordinary star, a pulsar will move in its own small orbit if it
has a planet thus giving rise to the Doppler effect. Calculations based on pulse-timing
measurements can then reveal the parameters of that orbit.
Four planets have been detected with this method so far. The three initially discovered
pulsar planets (around PSR1257+12) have masses of 0.02, 4.3, and 3.9 Earth masses
with the respective orbital periods of 25, 66 and 98 days.
The formation of pulsar planets is still discussed. It could be planets that survived the
supernova that created the pulsar. More likely, the planets formed after the supernova
from remaining debris or “fallback”. The latter model is supported by recent
observations with the Spitzer Space Telescope that identified a disk of debris around a
pulsar. Pulsar planets would be entirely incapable of supporting any form of life as we
know it due to the colossal amounts of electromagnetic radiation emitted by pulsars.
Figure 10: Period variations of PSR1257+12. Each period measurement is based on observations made
on at least two consecutive days. The solid line denotes changes in period predicted by a two-planet
model. Later measurements have actually identified a third, smaller planet in the same system. From
Wolszczan & Frail, 1992.
4.2 Properties of extrasolar planets
In 1992 the detection of the first extrasolar planets (namely pulsar planets) has been
announced (Wolszczan & Frail, 1992). The discovery of the first extrasolar planet (51
Peg b) around a solar-type star followed in 1995 (Mayor & Queloz, 1995). Until now
the number of detected extrasolar planets has increased to more than 1000 with masses
ranging from super-Earth (even one sub-Earth mass planet when accounting for pulsar
planets) to several times the mass of Jupiter. This demonstrates that planet formation is
not an extraordinary event but a common occurrence.
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Our solar system forms the basis for most of our information about how planetary
systems must develop. However, the degree to which it is actually representative of all
planetary systems is unclear. It is certainly very different from all extrasolar planetary
systems discovered so far. But it should be kept in mind that the observations of
extrasolar planets are still strongly biased. In this section we will discuss the basic
statistical properties of the extrasolar planets such as the period, size, and eccentricity of
the orbit, the mass, and properties of host stars. These quantities are relatively easy to
obtain for the bulk of all known extrasolar planets. Recently, exoplanetry research has
entered into a second phase by achieving direct observations of exoplanets for a few
selected cases. We will look at these exciting new developments in Chapter 3 when we
discuss planetary atmospheres.
Up to date information on detection of exoplanets:
The Extrasolar Planets Encyclopaedia (http://exoplanet.eu/)
Orbital radius of extrasolar planets:
Some stars have giant planets orbiting at distances up to 10 times closer to their star
than Mercury to the Sun (Fig. 11). While not all are that close, a significant number of
them orbit within 0.1 AU of their star! This was very surprising as the temperatures
in the protoplanetary disk this close to the central star is expected somewhere around
2000 K, far too hot for the existence of small solid particles. Planets around 0.04-0.05
AU have orbital periods of 3-4 days and are referred to as hot Jupiters.
Figure 11: Mass vs. semi-major axis of extrasolar
planets (red trangles) and solar system planets
(green squares). The fact that only high mass
exoplanets have been detected is of course a
observational bias and due to the sensitivity of our
instruments.
(From http://jilawww.colorado.edu/~pja/)
The presence of these giant planets at close orbital distances requires significant
modifications and/or extensions to the standard formation model for three major
reasons. First, the lack of small solid particles due to the high temperature within
0.1 AU makes it very difficult, if not impossible, for the standard core accretion model
to build up planets. Second, the mass of a typical protoplanetary disk within the orbit
of the closest objects observed would not amount to a Jupiter mass by a large factor
even assuming 100% efficiency in collecting the matter. Third, even if there was
sufficient mass available, the young 51 Peg b (first discovered extrasolar planet) for
example would be torn apart during a formation at its current location by the star's
gravitational forces.
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Eccentricities of orbits of extrasolar planets:
Except for the very close planets for which tides circularize the orbit, the eccentricity of
many extrasolar planets is rather large. To illustrate to what extend these systems
differ from our own solar system, we have plotted the eccentricity of k nown extrasolar
giant planets as a function of their semi-major axis (Fig. 12) and as a function of their
orbital period (Fig. 13).
Interestingly, according to Fig. 13 there is a clear pile-up of planetary companions with
periods around 3 days combined with an apparent absence of planets with shorter
periods. This is in complete contrast to the period distribution of stellar companions,
which can have periods much shorter than 3 days. Meanwhile, since the compilation of
Fig. 13, a few giants with periods as short as 1.2 days have been detected, but the
overall picture has not been altered.
These results not only indicate a different formation mechanism of stars and planets but
also imply that the processes involved in the planetary migration makes the planet
“stop” at a distance corresponding to about 3 days.
Figure 12: Eccentricity as a function of semimajor axis for giant extrasolar planets (red
triangles) and the solar system planets (green
squares). Note the difference in orbital
parameters of giant planets in and outside our
own planetary system. Extrasolar planets have
in general much larger eccentricities.
Surprisingly a large number of giants exist very
close to the central star. The absence of
extrasolar planets beyond 3 AU is due to
observational biases. See more recent plots at
http://exoplanetarchive.ipac.caltech.edu/
Figure 13: Eccentricity as a function of
orbital period for extrasolar planets (red open
pentagons), for binary stars (filled circles), for
the solar system giants (green stars in lower
right corner), and for Earth (usual symbol in
blue). At first glance extrasolar planets and
stellar binaries have similar orbital
parameters, However, there exists a clear
statistical overabundance of planets with a 3
day period. This is in sharp contrast with the
period distribution of stellar companions,
despite the fact that a few planets with periods
as short as 1.2 days have been detected since
this figure has been compiled.
From Santos et al. (2002).
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Masses of extrasolar planets:
Very important information is brought to us by the analysis of the mass spectrum of the
planetary companions. It is expected that the different type of formation of stars
(collapse) and planets (core-accretion), which could be understood as two different
classes of objects, might be visible in the mass distribution of low-mass stellar
companions. In fact this information could be employed to refine planet formation
models. However, observations of low-mass companions of solar type stars are still
very biased and have to be interpreted with care.
Figure 14 shows such an example of biased observations. It gives the mass distribution
of companions to solar-type stars, indicating a clear discontinuity for the mass regime
between about 20 and 60 times the mass of Jupiter: there are basically no
companions found (as of 2002!) having those masses (note that the mass of Jupiter is
103 M). This result appeared even more striking if we note that the observational
technique used so far to search for extrasolar planets is more sensitive to massive
companions than to their lower mass counterparts.
Figure 14: Distribution of minimum masses for the
currently discovered low-mass companions to solar-type
stars. A clear gap is visible in the range of 20–60 Jupiter
masses, the “Brown Dwarf desert”, supporting the view
that different mechanisms are involved in the formation
of stars and planets. Note, however, that newest data
since 2006 starts to fill up the gap with Y dwarfs. From
Santos et al. (2002).
This gap, usually called the “Brown Dwarf
desert”, separates the low mass “planetary” companions from their high mass “stellar”
counterparts, and was believed to tell something very important about the physical
processes involved in the formation of these two populations, namely that stars, even
the low mass ones, are thought to be formed as the result of the gravitational collapse
and fragmentation of a cloud of gas and dust, while a planet forms in a circumstellar
accretion disk. In this sense these results seemed to fit perfectly into the theory.
However, newest observations of the past year start to fill up this gap with a new class
of sub-stellar objects called Y dwarfs, which are a subclass of the brown dwarfs. Brown
dwarfs are stars not massive enough that hydrogen burning can be initiated in the center,
i.e. less than about 80 Jupiter masses (MJ). Brown dwarfs heavier than 13 MJ do fuse
deuterium and above roughly 65 MJ fuse both deuterium and lithium.
Since the Brown Dwarf desert starts to be filled up, the question of how to distinguish
brown dwarfs and giant planets arises, i.e. how these objects should be exactly defined.
One possibility is to distinguish by formation mechanism: planets form by accretion,
brown dwarf and stars by collapse. Another possibility is that brown dwarfs are required
to have experienced fusion at some point in their history, which would draw a clear
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maximum mass of 13 MJ for planets. However, there is evidence that hot Jupiters (e.g.
HD 209458 b) evaporate because they are so close to the central star. The same could
maybe happen to Y-dwarfs so they become hot Jupiters.
In principle this story is just a question of definition. We will see how it will be decided
in the future. From the physical point of view it is however of great interest to learn
more about these objects and their formation history.
Metallicity of host stars:
Host stars have, on average, a higher metal content than stars with no planetary
companions detected. In other words, these stars have a higher ratio of heavy elements
to hydrogen than that observed in average solar-type field stars. More than 20% of stars
with metallicity greater than two times the solar metallicity harbor a planet, whereas
only 3% of stars with solar metallicity have a giant planet (Fig. 15). However, this
does not imply that giant planets cannot be formed around metal-poor stars. Rather, it
suggests that the probability of formation in
such a case is substantially lower.
Figure 15: Percentage of stars that were found to have
planets among the Geneva planet search survey sample
as a function of the relative amount of iron (i.e.,
metallicity) with respect to the Sun. This figure shows
that about 25% of the stars with twice the solar
metallicity harbor a planetary mass companion,
whereas this percentage decreases to below 5% for
stars with the same metal content as our Sun. From
Santos et al. (2005).
A possible and likely interpretation of this
may be that the higher the metallicity of the cloud that gives origin to the star/planetary
system (and thus the higher the dust content of the disk), the faster a planetesimal can
grow, and the higher the probability that a giant planet is formed before the
protoplanetary disk dissipates. In other words, the metallicity seems to be playing a key
role in the formation of the currently discovered extrasolar planetary systems.
Internal structure:
In the case of a transiting planet we can determine the radius of the planet and the
orbital inclination (transit method). Combined with the radial velocity method we thus
find the true mass as well as the average mass density of the planet. Therefore, in
combination with some modeling, it becomes possible to infer information about the
internal structure of such planets (Fig. 16).
The range of mean densities identified in this way for hot Jupiters varies significantly.
Explaining an exoplanet with a similar mass to Jupiter but a smaller size (higher mean
density) is easy — it probably has a larger core, with more heavy elements in general.
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But for exoplanets that are less dense, there comes a point when even a planet made just
of hydrogen isn’t enough to explain its low density. Indeed the known low-density hot
Jupiters pose a theoretical challenge and probably require some additional internal
energy source not considered in current models. On the low density side, the most
extreme case found so far is TrES-4 that has the largest radius (1.7 RJupiter) and lowest
density (0.2 g/cm3) of any of the known transiting planets.
Figure 16: Cut-away diagrams of Jupiter, Saturn, and two extreme cases of extrasolar planets, drawn to
scale. The observed radius of HD 149026b implies a massive core of heavy elements that makes up
perhaps 70% of the planetary mass. In contrast, the radius of HD 209458b intimates a coreless structural
model, as well as an additional energy source to explain its large radius. From Charbonneau et al. (2007).
4.3 Implications on planet formation theory
Planet migration:
To reconcile theory and observations different mechanisms have been considered which
essentially allow planets to migrate from their birth place to where they are
observed today. This planetary migration is not a new idea, but it was never considered
before as an essential ingredient in planet formation.
Migration can be due to several physical processes such as gravitational scattering in
multiple systems (possibly involved in the formation of Uranus and Neptune) or
gravitational interactions between the gaseous and/or the planetesimal disk and the
planet (that lead to density perturbations, tidal forces, and torques exerted on the
planet). These two mechanisms must necessarily occur, and interactions between an
embedded planet and a gaseous disk were discussed before the discovery of the first
extrasolar planet. The question is therefore not whether migration takes place or not but
rather what its direction and amplitude are.
For migration based on gravitational interactions two types of migration modes have
been identified, depending on whether the planet is massive enough to open a gap in the
disk (type II migration) or not (type I migration) (Lin et al. 1996, Ward 1997). Typically
a planet would start with type I migration, which is relatively fast due to friction and
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gravitational interaction with the viscous disc. Because the disk material rotates faster
inside the planets orbit than outside, the disk exerts a torque on the planet, which results
in a loss of angular momentum and an inward spiraling planet. If a planet becomes
sufficiently large to accrete gas it will cause a gap to form in the disk on the path of the
planets orbit. This slows down migration drastically by a factor of 10 to 100. The
inward drift continues as type II migration at the same rate as the inward drift of the
disk material (due to accretion onto the central star) occurs. All these migration models
conclude that planets are migrating mostly inward toward the star, over large
distances and relatively fast. In fact, migration time scales (a few 105 to a few 106 years)
obtained so far are so short (especially for type I migration) that, in almost all cases,
planets should not survive but should fall into their host star.
This raises the question whether mechanisms exist for stopping inward migration
(otherwise, why would we observe so many planets?). Evidence for such a mechanism
may be deduced from the observed overabundance of systems with periods around 3
days (Fig. 13). Physical mechanisms responsible for halting and parking a planet at
short distances from the host star include the existence of a central cavity in the disk
(cleared by the magnetosphere of the central star) or tidal friction (angular momentum
exchange between the planet’s orbital motion and the spin of the star).
Although these stopping mechanisms are relevant at short distances, they do not explain
why giant planets are found at intermediate distances (e.g., with periods around 1
year) nor why Jupiter, for example, has apparently remained beyond 5 AU. In fact, the
recent addition of disk evolution and planetary migration mechanisms into the coreaccretion models suggests that planets essentially migrate until the disk disappears (in
fact, until the disk becomes much less massive than the planet).
Disk instability model
Boss (1997, 2003) has suggested that besides the core accretion scenario, giant planets
might also be formed as a result of disk instability processes. In this scenario, giant
planets form directly from the gravitational fragmentation and collapse of the
protoplanetary disk within a few dynamical time scales, i.e. in about 103 years (Fig. 17).
The short time scale of planet formation in the disk instability model represents one
of its biggest advantages over the core accretion model.
However, it has its own drawbacks. For example, with the direct collapse scenario of
the disk instability model it is more difficult to explain the rocky and icy cores of the
giants (in the solar system; for extrasolar planets we still know little about the presence
of cores, but see discussion to Fig. 16 above). Further, disk instability is not very
dependent on the metallicity. In other words, if the disk-instability models were the
most important mechanism involved in the formation of giant planets, we should not
expect to see a strong dependence on the rate of planet detection as a function of the
metallicity. The huge dependence observed is thus probably a sign that the core
accretion scenario is the important mechanism involved in the formation of giant
planets. But remember that observations of extrasolar planets might still be biased.
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The two formation paradigms are currently being critically examined. The core
accretion model is sufficiently advanced to begin to allow quantitative calculations to be
made and thus permits a direct comparison with giant planets in and outside our solar
system. The direct collapse model is in a state where only qualitative statements can be
made without the possibility to compare quantitatively with observations.
Figure 17: Illustration of the disk instability model for forming giant planets. These face-on snapshots
from 3D simulation show the density (increasing from dark blue to magenta, red and yellow) of a
protoplanetary disk at 160 years (left) and 350 years (right) of evolution. After about 150 years the disk
develops trailing spiral arms. A two-armed mode grows up to the point, where after about 200 years,
fragmentation occurs along the arms, and more than one distinct clump appears. From Quinn (2003)
4.4 Projects and Facilities
The detection of 51 Peg b in 1995 has opened a new race to search for and study
extrasolar planets. Since then many instruments have been dedicated more or less to this
field and many new observatories are planed for the future.
The first goal is to collect better statistics on basic parameters, such as orbital
parameters, properties (e.g. metallicity) of host stars, and the probability that stars have
a planet or a planetary system. This allows us already to refine the planet formation
theory. In the second phase (which has already started for a few selected planets), with
better instruments and as soon as direct observations are achieved, we can study
physical and chemical properties of the planets, such as temperature, abundances, and
surface structure. At the same time a lot of effort is invested into detecting planets with
longer orbital periods and smaller planets, in particular terrestrial, Earth-like planets.
In the case of giants we are naturally ahead. Collection of statistics of basic parameters
is ongoing. By now Neptunes, super-Earths and planetary orbits with larger semi-major
axis can be detected. In fact, first steps in studying the physics of planets (composition,
temperature distribution) are undertaken, thanks to observed transits and infrared
observations with the Spitzer Space Telescope.
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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The observation of Earth-like
terrestrial planets is obviously more
challenging.
Nonetheless,
the
collection of basic statistics will
become feasible within the next 5 to
10 years. Subsequent refinement
and direct observation of light from
terrestrial planets is then just a
question of time, possible within
another 10 years. In the case of
terrestrial planets we want to go one
step further than in the case of
giants. The ultimate science goal is
the detection of signatures of life on
extrasolar planets.
In the following we have a look at
the most exciting projects, either
ongoing or in the planning phase,
with an emphasis on space projects.
It should however be understood
that many more observatories, in
particular also ground based, will be
heavily employed in the search of
extrasolar planets.
Figure 18: Detection capabilities of current and future
observatories.
Spitzer Space Telescope (NASA)
Basic facts:
 Launch:
25 August 2003
 Orbit:
heliocentric, Earth trailing
 Web page: http://www.spitzer.caltech.edu/
Instruments:
 85 cm telescope
 Infrared imaging, infrared spectroscopy, and infrared spectrophotometry
Method for studying planets:
 Transit method (primary and secondary eclipses)
Science:
 Ideal for cold objects, due to observations in the infrared
 Protoplanetary and circumstellar disks, dust, molecular clouds
 Extrasolar planets (hot Jupiters, transit method, thermal emission of planet,
temperature distribution day/night side)
 Early Universe (distances where main light emission of galaxies is redshifted
into infrared!)
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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Keck Interferometer
Basic facts:
 Ground-based, Mauna Kea, Hawaii
 Operational since 2002
 Web page:
http://planetquest.jpl.nasa.gov/Keck/keck_index.cfm
Instruments:
 Two 10-meter telescopes
 85 meter baseline interferometer
 Nulling interferometry (because only two telescopes, cancels light from a star so
that faint dust surrounding star can be observed; interferometric imaging not
possible)
 Near-infrared (K waveband)
Method for studying planets:
 Direct imaging (spectroscopy) and astrometry
Science (within field of extrasolar planets):
 Survey protoplanetary disks, dust around nearby stars
 Spectroscopy of hot Jupiters and astrometry of giant planets
Large Binocular Telescope Interferometer (LBTI)
Basic facts:
 Ground-based, Mount Graham, Arizona
 Since ~2012
 Web page:
http://lbti.as.arizona.edu/
http://planetquest.jpl.nasa.gov/lbti/lbti_index.cfm
Instruments:
 Two 8.4 meter mirrors
 15 meter baseline nulling interferometry
 Near-infrared (J, H, and K wavebands)
Method for studying planets:
 Direct imaging
Science (within field of extrasolar planets):
 Survey protoplanetary disks, dust around nearby stars
 Direct imaging of giant planets
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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COROT (“Convection Rotation and Planetary Transits”, ESA)
Basic facts:
 Launch:
2006
 Orbit:
Polar orbit around Earth
 Web page:
http://smsc.cnes.fr/COROT/
http://sci.esa.int/science-e/www/area/index.cfm?fareaid=39
Instruments:
 27 cm telescope
 optical
Method for studying planets:
 Transit method
Science:
 Transition region between giants and terrestrial planets: first detection and
orbital parameters
 Goal: ~10 large terrestrial planets (several times larger than Earth, closer than
0.5 AU to host star) and several hundred Jupiters
 Stellar seismology (to study internal structure of stars, more than 120’000 stars
surveyed)
Kepler (NASA)
Basic facts:
 Launch:
2009
 Orbit:
heliocentric, Earth trailing
 Web page: http://www.kepler.arc.nasa.gov/
Instruments:
 1.4 meter telescope
 optical
Method for studying planets:
 Transit method (100’000 targets observed simultaneously)
Science:
 Survey distant Earth-sized planets: orbital parameters, how common? Different
spectral classes of host stars.
 Identification of potential targets for later, more advanced missions
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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Sphere (“Spectro-Polarimetric High-contrast Exoplanet REsearch”, ESO)
Basic facts:
 Ground-based, VLT, Paranal, Chile
 Commisioned 2014
 Web page:
http://www.eso.org/projects/aot/vltpf/
Instruments:
 New instrument on VLT
(Very Large Telescope; four 8 meter telescopes with interferometry)
 Infrared imaging and spectroscopy
 3 science subsystems, including imaging polarimeter ZIMPOL (ETH Zurich)
Method for studying planets:
 Polariztion method
 Direct imaging
Science:
 Earth-like terrestrial planets
 Composition, thermal emission
GAIA (ESA)
Basic facts:
 Launch:
2014
 Orbit:
Lagrange point L2 of Sun-Earth system
 Web page: http://sci.esa.int/science-e/www/area/index.cfm?fareaid=26
Instruments:
 Dual telescope for astrometry
 Optical photometry
 Radial velocity spectrometer
 Spatial accuracy: 20  arc sec
Method for studying planets:
 Astrometry
 Radial velocity
Science:
 Catalogue 1 billion stars  3D-map of Milky Way
 Expected to find 10’000–50’000 Jupiters within 150 light years from Earth,
having periods up to 9 years
SIM Lite (“Space Interferometry Mission”, NASA)
Basic facts:
 Launch:
 Orbit:
not allocated (budget problems)
heliocentric, Earth trailing
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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
Web page:
http://planetquest.jpl.nasa.gov/SIM/sim_index.cfm
Instruments:
 9 meter basline interferometer in the visible
 Spatial accuracy: 1–4  arc sec
Method for studying planets:
 Astrometry
Science:
 Earth-sized terrestrial planets (at 1 AU) around nearby stars, basic parameters
Terrestrial Planet Finder Coronograph (TPF-C, NASA)
Basic facts:
undefined (originally 2016, funding problem)
 Launch:
 Orbit:
Lagrange point L2 of Sun-Earth system
 Web page:
http://planetquest.jpl.nasa.gov/TPF/tpf_index.cfm
Instruments:
 Coronograph (i.e. obscuring light from central star)
 optical
Method for studying planets:
 Direct Imaging
Science:
 Study Earth-like planets
 Chemistry
 Atmosphere
 Surface structure
 Biomarkers, signatures of life
Terrestrial Planet Finder Interferometer (TPF-I, NASA) / Darwin (ESA)
Basic facts:
 Launch:
undefined (originally 2020, funding problem)
 Orbit:
Lagrange point L2 of Sun-Earth system
 Web pages:
http://planetquest.jpl.nasa.gov/TPF/tpf_index.cfm
http://www.esa.int/esaSC/120382_index_0_m.html
 Possible joint project NASA/ESA
Instruments:
 Interferometry
Astrobiology: 4 Extrasolar planets: detection, properties, projects
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
Infrared
Method for studying planets:
 Direct Imaging
Science:
 Study Earth-like planets
 Chemistry
 Atmosphere
 Surface structure
 Biomarkers, signatures of life
References
Beaulieu, J.-P. et al. 2006, Nature 439, 437
Boss, A. P. 1997, Science, 276, 1836
Boss, A. P. 2003, ApJ, 599, 577
Charbonneau, D. et al. 2007, in “Protostars and Planets V”, eds. B. Reipurth, D. Jewitt,
& K. Keil, University of Arizona Press, 701
Lin, D. N. C., Bodenheimer, P., Richardson, D. C. 1996, Nature, 380, 606
Mayor, M., Queloz, D. 1995, Nature, 378, 355
Quinn, T. 2003, How to Cook a Giant Planet, Projects in Scientific Computing,
http://www.psc.edu/science/
Santos, N. C. Mayor, M., Queloz, D., Urdí, S. 2002, Extrasolar Planets, The ESO
Messenger, 110, 32
Santos, N. C., Benz, W., Mayor, M. 2005, Science, 310, 251
Ward, W. 1997, ApJ, 482, L211
Wolszczan, A., Frail, D. A. 1992, Nature, 355, 145
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