Figure 1. The
innermost regions of the protoplanetary disk surrounding TW Hydrae are unveiled
in this composite of images captured by the ALMA instrument. TW Hydrae is a newborn Sun-like star located 176 light
years away. Sean Andrews & al. (2016) propose that the gap and ring
structure visible in the inset image has been carved by a rocky planet
accreting at a circumstellar radius of about 1 astronomical unit (1 AU),
equivalent to the distance of Earth from the Sun. However, Barbara Ercolano &
al. (2016) argue that the dust-free gap results from photoevaporation of the
inner disk by X-ray flux from the central star. Both studies suggest that the
gaps in the outer disk represent the condensation fronts of various chemical
species, such as carbon monoxide and molecular nitrogen. Current theory
identifies these structures as potential planet
traps.
------------------------------
Despite recent interest in the theory of in
situ assembly, most astronomers
over the past two decades have relied on models involving planet migration to
explain the origin of planetary systems. These models are directly informed by
our understanding of the primordial nebulae that surround newborn stars. The previous
installment of this three-part series offered a basic picture of protoplanetary disks (the preferred
term for planet-forming nebulae) and reviewed the pros and cons of in situ formation. This installment
explores the recent explosion of theoretical studies that invoke gas-driven
migration in radially structured protoplanetary disks as the principal mechanism
underlying planet formation and system architecture.
Intensive discussion of radially structured
disks dates back at least 10 years, to the 2006 publication of a study led by
Frederic Masset: “Disk surface density transitions as protoplanet traps.” Since
then, the term planet trap has
appeared regularly in articles on system evolution, with a currency even wider
than the extended circle of researchers who collaborate with Masset. Such traps
are a key example of the structures invoked in this emerging field.
A major shortcoming of existing migration models
provided the jumping-off point for the current approach: growing recognition that
the combined effects of aerodynamic drag and gas-driven migration, as modeled
in theory, would likely sweep all solid particles in the nebula into the
central star before they had time to form protoplanets.
Without those building blocks, planetary systems could never evolve.
Yet obviously they do. Accordingly, theorists
have searched for some mechanism that could stall migration and enable solid
mass to accumulate at various radial locations. These efforts have coincided
with increasingly precise measurement and imaging of nearby protoplanetary
disks, which are beginning to offer evidence of disk structures that are
consistent with planet traps (Figure 1).
So this posting will look at 1) the
basic elements of migration theory, 2) a sample of the disk structures proposed
by current scenarios, 3) condensation
fronts and their evolution, 4) the formation of the first planetesimals, 5)
the late stages of disk evolution, and 6) the nativity and composition of the
known exoplanets.
Figure 2. Protoplanetary disk model
Schematic view of a protoplanetary disk
surrounding a Sun-like star, seen edge-on. The large lavender shape represents the
gas nebula, which is made of hydrogen & helium in an approximate ratio of
3:1. The vertical profile of the gas flares in the outer region of the disk,
beyond 10 AU. Despite a gas-free central cavity inside 0.05 AU, gas molecules
flow continually through the disk to the star and deposit mass through the
stellar magnetic poles. Suspended throughout the disk are solid particles of
rock and frozen volatiles, the building blocks of planets. Rocky particles
dominate the hot inner disk; icy particles dominate the cold outer disk. As the
disk evolves, solids settle at its midplane, which is likely aligned with the
star’s rotational axis. Accumulation of solids favors the assembly of pebbles
and planetesimals, which collide to form protoplanets (also known as embryos or
cores).
------------------------------
1) migration modes
In a protoplanetary disk, nothing stands
still. Gas molecules orbit the central star and dust grains are carried along
with the flow of hydrogen and helium (H/He) as it streams into the gravity well
created by the star’s mass (Figure 2).
Current models tell us that the process of accretion is driven by turbulence
within the nebula. The calmest region of the gas disk is the midplane, where dust
settles and aggregates, encouraging collisions. Micron-sized grains are
lightweight enough to travel at the same speed as the gas, but aggregations
measurable in centimeters – that is, pebbles
– experience aerodynamic drag that causes their orbits to decay.
The problem of orbital decay gets worse as
aggregates grow. When a protoplanet approaches the mass of Mars (~0.1 Earth
masses or Mea), it becomes subject to Type
I migration. Objects ranging up to and beyond the mass of Neptune are
similarly affected. Type I migration is caused by the interaction between the
flow of gases and a forming protoplanet, such that the object is simultaneously
subject to positive and negative torques. When the negative torque exceeds the
positive torque, as is usually the case, the object migrates inward. When the
reverse is true, the object migrates outward. In some models, Type I migration
can deliver an Earth-mass protoplanet originally traveling on an Earth-like
orbit to the threshold of its parent star in about 100,000 years (Baillie et
al. 2015). More massive objects travel even faster, such that on object of 10
Mea originally orbiting at 5 astronomical units (AU) can migrate to the inner
edge of the disk in less than 30,000 years (Chambers 2006).
Figure 3. Type II migration of a baby gas giant
The blob in the dark ring is the forming
planet, which is being carried inward by the flow of accretion as it orbits in
a dust-free lane. Spiral streams of gas link the planet to the inner and outer
disk. (Source: Figure 32 of Armitage 2010)
------------------------------
Still larger objects can escape this process.
Depending on the vertical extent of the gas disk at the site of formation, a
protoplanet approaching the mass of a gas giant (e.g., 50-150 Mea) can open a
gap in the disk, meaning that it clears all the gas from its vicinity and
orbits alone in an empty lane around the star. The exact mass required for gap clearing
depends on the scale height of the disk, which changes over time. Disks are at
their thickest and most massive in extreme youth, attenuating incrementally as
they age. Only a massive object can open a gap in a thick disk, whereas a
protoplanet as lightweight as Neptune might be able to open a gap in a
sufficiently attenuated disk. Once ensconced in its gap, the planet undergoes Type II migration (Figure 3), traveling inward with the flow of gas at a more
leisurely rate than in the Type I regime. Nevertheless, it still migrates, so
that the orbit of a young gas giant in this mode might shrink from 5 AU to 0.1
AU in half a million years (Chambers 2006).
2) barriers and traps
For as long as the various modes of drag,
drift, and migration have been understood, astronomers have realized that some
opposing mechanism must be available to ensure that pebbles don’t fall into
stars before they accrete into protoplanets, and that protoplanets don’t get
stranded at the inner edges of disks before they can establish cooler orbits. In
recent years, theorists seem to be converging on a scenario in which the gas
dynamics in specific regions of a protoplanetary disk counterbalances negative
torque with positive torque, for a net of zero. In this way, solid particles can
remain stranded in one place long enough to accrete into planetesimals,
protoplanets, and planets. Such regions have been identified as the sites of “transitions”
(Masset et al. 2006), “inhomogeneities” (Hasegawa & Pudritz 2011), or “irregularities”
(Baillie et al. 2016).
Several candidates have been proposed as the
agents of this trapping or focusing process. Collectively, they are known as
planet traps. Among the most consistent choices are condensation fronts (or
sublimation lines) in the evolving disk – locations where specific chemical
species condense (or sublime) out of the disk gases (Figures 4 & 5). The most widely discussed of these fronts is
the ice line or snow line: the
minimum radius at which water freezes instead of sublimating. Another is the silicate condensation line, where
silicate dust vaporizes, creating an inner edge for the solid component of the
disk (Morbidelli et al. 2016). In many models, this edge casts a long shadow on
the midplane of the disk, contributing to a reduction in temperature.
Condensation fronts act as planet traps
because they create a bump in the local surface density of solids (Masset et
al. 2006). However, other candidates for this role are available. In a series
of publications, Yasuhiro Hasegawa and collaborators have detailed a total of
three potential planet traps (2010, 2011, 2014). In addition to condensation fronts – to which they
apply the generic term ice line, regardless of chemical species – they include dead zones and heat transitions.
Many theoretical treatments of planet traps
specify that they are intended to apply only in cases of Type I migration, with
some even defining a relatively narrow range of masses for affected objects.
Others, however, discuss structures capable of capturing a broad distribution
of masses, including gas giants undergoing Type II migration.
Figure 4. Water condensation front around V883 Orionis
This image by the ALMA instrument reveals a protoplanetary disk of about 0.3 Msol with a
radius of about 125 AU surrounding the protostar V883 Orionis, which masses about 1.3 Msol and resides 415 parsecs (1350
light years) away in the Orion Nebula Cluster (Cieza & al. 2016). The dark
ring corresponds to the system ice line or water
condensation front, currently located at a circumstellar distance of about
42 AU. Normally the ice line would be much closer to a star of this age and
mass (in the range of 5-10 AU), but an abrupt accretion event in which the star
ingested a substantial amount of hydrogen from the disk caused an outburst. The result was a temperature
spike throughout the disk that evaporated water inside a radius similar to the
orbit of Pluto in our Solar System. Image credit: ALMA (ESO/NAOJ/NRAO) / L.
Cieza.
-----------------------------
2a) dead zones
A dead zone is a region in the nebula where
turbulence and viscosity are suppressed and the inward flow of gases slows down.
This concept has been discussed for at least 20 years as the possible source of
several phenomena associated with protoplanetary disks. (I have to wonder if
its coinage was related to Stephen King’s 1979 horror novel, The Dead Zone.) Dead
zones have been proposed to explain outbursts like those observed in newborn
stars such as FU Orionis, where
the sudden accretion of a large quantity of trapped mass stimulates a burst of stellar
radiation (Gammie 1996). They have been described as a mechanism for trapping
dust grains at specific locations in the disk, promoting the rapid collapse of
solids into planetary embryos the size of Mars (Lyra et al. 2008). They have even
been anointed as the saviors of planetary systems through their ability to halt
Type II migration, which would otherwise deliver forming gas giants to
star-hugging orbits (Matsumura et al. 2007).
In the classic model, the magnetic field of a
newborn star induces magnetorotational instability in the protoplanetary
nebula. This instability creates turbulence, which drives the flow of H/ He
into the central star. Although the existence of magnetorotational instability
has recently been questioned (Morbidelli & Raymond 2016), theorists still
seem to agree that turbulence is real, regardless of the source. Current models
still indicate that its suppression – by whatever means – will reduce viscosity
and halt inward migration.
Some discussions appear to confine the dead
zone to the midplane of the protoplanetary disk, without emphasizing its radial
location (Martin & Livio 2012). However, Hasegawa and Pudritz (2011)
describe their dead zone as a high-density region confined to the inner disk,
where turbulence is calmed and the accumulation of dust at the midplane is
enhanced. They note that an inner dead zone can coexist with other kinds of
planet traps, including ice lines (which they also describe as opacity
transitions) and the heat transition. They even characterize ice lines as a
species of “self-regulated, localized dead zone.”
The heat
transition proposed by these authors and others refers to the radius in a
protoplanetary disk where the source of heating shifts from friction caused by
the flow of gases (viscous heating) to energy radiated by the central star
(irradiative heating). Models of disks around Sun-like stars place this transition
around 10-12 AU (Baillie et al. 2016), where the disk begins to flare. Flaring
enables the outer disk to escape the shadow of the inner dust wall, so that surface
gases heat up and ionize. At the midplane, however, icy solids still accumulate.
Hasegawa & Pudritz (2011) argue that the placement of
the various planet traps defines the initial
orbital distribution of protoplanets. Although all traps tend to move inward
over time as the rate of gas flow through the disk declines, each type of trap moves
differently, resulting in different outcomes for its complement of protoplanets.
Hasegawa & Pudritz also argue
that the positioning of planet traps depends strongly on the mass of the
central star and its rate of gas accretion. Through this dependency, they
conclude, “host stars establish preferred scales [for] their planetary systems.”
Figure 5. Carbon monoxide condensation front around TW
Hydrae
A
doughnut made of carbon monoxide surrounds TW Hydrae, a young Sun-like star, as
captured in this image by the ALMA instrument. Qi & al. (2013)
estimate its inner radius at 30 AU. (Note that the yellow dot in the center of
this figure is an artificial marker of the position of the host star – it is
not present in the original image.)
------------------------------
2b) condensation fronts
Kevin Baillie & colleagues (2016) follow
Hasegawa & Pudritz in emphasizing the role of condensation fronts (which
they call sublimation lines) as well as the heat transition barrier. Although
their terminology foregrounds “lines,” they explain that each line actually marks
the center of a broad zone or plateau that extends over a radial distance of about
1 AU. In addition to water and silicates, they model the sublimation of
volatile and refractory organic molecules, along with troilite, an iron sulfide
mineral. These three types of molecules sublime at temperatures between those
of silicates and water, creating a series of tightly spaced traps between 0.1
and 1 AU. However, the authors find that the traps created by troilite and
volatile organics are less effective than the others, while the silicate
condensation trap is less enduring than the water trap.
Like other theorists (Pollack et al. 1996,
Masset et al. 2006, Hasegawa & Hirashita 2014, Bitsch et al. 2015), Baillie
& colleagues also foreground the effects of opacity on disk evolution. Local
opacity is determined by the size of ambient dust grains, with micron-sized
particles increasing opacity and millimeter to centimeter sizes (pebbles) reducing
it. They note that condensation fronts are associated with opacity transitions, since solids accrete more readily outside these
fronts than inside them, causing an increase in opacity. Again, however,
troilite and volatile organics make a smaller contribution to opacity than do
water and silicates.
While Baillie & colleagues do not
discount the potential of dead zones for trapping planets, they have not
explicitly integrated this type of structure in their extant models. Instead,
they focus on the behavior of migrating objects that arrive at planet traps,
paying special attention to mass. They argue that condensation fronts can trap
low-mass planets, but not gas giants, for which the heat transition is the only
effective barrier. They also find that the condensation fronts and heat
transition zone create planet deserts
about 1 AU beyond the inner edge of each structure. Thus, each plateau is
bounded by an outer wasteland where no worlds can grow.
2c) outward migration
Another emphasis of their work is the
existence of zones of outward migration associated with the heat transition,
the silicate condensation front, and the water condensation front. These
structures can confer a positive torque on low-mass planets, adding further
complexity to an evolving system’s orbital architecture. Nevertheless, all
planet traps move inward over time with the cooling and attenuation of disk
gases. Thus, even if a forming planet is stalled in a trap, its orbit will
still shrink as the disk ages.
Like Baillie’s group, Bitsch & colleagues
(2016) omit dead zones from their modeling to focus on opacity transitions
associated with thermal discontinuities, disk geometry, and local surface
densities of solids. Similarly, they explore the zones of outward migration
created by temperature changes. Unlike Baillie’s group, however, they find that
more massive planets are more likely to migrate outward than low-mass planets.
They are particularly interested in metallicity,
which they propose as a key determinant of the temperature profile of a
protoplanetary disk. They find that higher metallicity is associated with wider
zones of outward migration in more distant regions of the disk. These factors
can help to explain the well-known correlation between stellar metallicity and
the presence of gas giant planets. (See Cossou et al. 2014 for an analogous treatment
of metallicity and migration.)
2d) giants at the gate
One final, widely discussed barrier to
drifting dust and migrating planets is the presence of a growing gas giant,
which depletes the gas and dust from its vicinity and opens a gap in the
nebula. A migrating planet that approaches such a radial gap will either stall
or be scattered, usually into a wider orbit. Izidoro & colleagues (2015) argue
that this behavior can explain the orbital architecture of our Solar System,
where the accretion of Jupiter outside the water condensation front created a
barrier to the inward migration of Saturn, Uranus, and Neptune. That may well
be so: the emergence of one or more gas giants might effectively segment the
disk into inner and outer domains. But the authors go on to predict an
anti-correlation between systems with gas giants and those with compact collections
of warm gas dwarfs. The architectures of a half-dozen systems – HD 10180, HD
219134, Kepler-48, Kepler-90, Kepler-167, and WASP-47 – demonstrate that this
prediction cannot be universally valid (see Almost
Jupiter).
Benitez-Llambay & colleagues (2015) use
the example of Jupiter to frame a different argument about the migration of
growing giants. They consider the case of a protoplanet of 3 Mea that has
already formed beyond the water condensation front. This hypothetical object
could easily stand in for baby Jupiter. Over a period of 100,000 years, it
doubles in mass by accreting a continual flux of colliding bodies. At the same
time, it responds to a negative torque that tends to shrink its orbit. But the
high temperatures produced by collisional accretion induce a heating torque that “constitutes a
robust trap against inward migration.” Depending on various factors, the
heating torque can cause the forming planet to decelerate its inward trajectory,
stall completely, or migrate outward. More massive planets are more likely to
experience outward migration, with a correlation between the object’s mass and
the radial distance it travels. Like Cossou’s group, Benitez-Llambay &
colleagues find that outward migration and full-blown gas giants are also more
likely to result from disks enhanced in metals.
Figure 6. Potential planet traps in
a young protoplanetary disk
Several sources of structure have been
proposed for evolving disks. A series of condensation
fronts (or snow lines or sublimation zones) arises in regions where local
temperatures cause common chemical species to condense. These molecules include
silicates, water, and carbon monoxide. In addition, the transition from viscous
heating in the inner disk to irradiative heating in the outer disk creates a
discontinuity known as the heat
transition. The radial location of these structures changes over time,
moving from larger to smaller radii along with a general cooling trend
throughout the disk.
-------------------
3) condensation fronts in space & time
Again and again, theoretical models converge
on condensation fronts as a primary source of disk structure as well as
exemplary planet traps. Many different chemical species have been nominated for
leading roles, with three standing out: silicates, water, and carbon monoxide.
Each one condenses around a specific temperature: silicates below 1500 K, water
below 170 K, and carbon monoxide below 70 K. However, the exact location in the
disk where each temperature obtains will vary widely, not only from system to
system but also from time to time within the same disk. This is because local
temperatures depend on irradiation, viscosity, and opacity. All these factors
vary over time, such that the location of each condensation front changes substantially
during disk evolution. Although the general trend is for each front to move
inward as the disk attenuates and cools, stellar flares and outbursts can cause
abrupt increases in temperature that evaporate dust grains and temporarily shift
all fronts outward (Cieza et al. 2016; see Figure
4).
Baillie & colleagues (2015) emphasize
that condensation fronts are regions rather than sharply defined lines,
preferring to describe them as plateaux. In addition, even their silicate condensation front is
relatively distant from the star during the earliest evolutionary phases.
Baillie’s group identifies an initial plateau extending from 0.3 to 1.3 AU, which
contracts to about 0.05-0.1 AU after “a few million years.” Similarly, Morbidelli
& colleagues (2016) propose that the inner edge of the silicate front in
the early Solar System was located at 0.7 AU, near the present orbit of Venus. Given
the extreme temperatures prevailing in the vicinity of the star, no solids can
survive in the inner disk and no accretion can happen there until the silicate front
shrinks substantially. Presumably this situation would complicate the models of
in situ accretion discussed in the previous
installment of this series, especially for exoplanets orbiting inside ~0.1-0.3
AU. It is conceivable that most objects on orbits shorter than about 100 days (i.e.,
the vast majority of Kepler planets) formed outside the original front and
arrived at their present location by inward migration after the inner system
cooled.
The plateau in surface density and pressure
created by the silicate front appears to coincide or at least overlap with the
dead zones proposed by Hasegawa, Lyra, and others. However, I haven’t noticed
any published treatments of their resemblance or lack thereof.
The water
condensation front is by far the most extensively discussed of any radial
structure, typically under the name of the snow line (or recently, snow region).
Since water is abundant in protoplanetary disks around Sun-like stars, the
local surface density of solids increases by a factor of 2 to 4 on orbits
outside the water condensation front (Kenyon & Kennedy 2008, Martin &
Livio 2012). This enhancement increases the likelihood that planetesimals and
protoplanets will form in this region.
The vagaries and excursions of the snow line have
been traced over the past few decades. Older literature typically defined 2.7
AU as the classical snow line in our Solar System, given the sharp transition in
the composition of small bodies observed at that location in the contemporary
Asteroid Belt. Purely rocky objects, like Vesta (semimajor axis = 2.36 AU),
orbit inside that radius, while objects that include a substantial proportion
of water ice, like Ceres (semimajor axis = 2.77 AU), orbit outside it.
However, more recent research argues that
this transition is simply the fossilized imprint of the position of the snow
line during the first few million years of system evolution, since the
present-day snow line is actually located inside Earth’s orbit (Ida &
Guillot 2016, Morbidelli et al. 2016). Thus, Kennedy & Kenyon (2008) have proposed
that the snow line migrates from about 6 AU to 1 AU during the first few
million years of disk evolution around a G-type star. Using a different
approach emphasizing dust opacity, Mulders & colleagues (2016) recently argued
for a primordial dispersion in the location of the snow line, with values
ranging from 1.4 AU to 8 AU for a G-type star. All these studies concur that
the snow line moves inward over time, and that its radial location at early
times leaves an indelible imprint on planetesimal accretion and ultimately
system architecture at later times.
The carbon
monoxide condensation front has received growing attention in the past few
years, especially since Qi & colleagues (2013) reported the detection of
this structure around TW Hydrae with the ALMA instrument (Figure 5). Since carbon monoxide freezes at much lower temperatures
than water, this condensation front is located much farther from a Sun-like
star than the other two major fronts. Qi & colleagues propose a radial
location of about 30 AU for stars of Solar mass, making the carbon monoxide
region the most accessible to direct observation. This structure might be
implicated in the formation of planets on very wide orbits (Dodson-Robinson et
al. 2008), and its effects might be especially significant for disks with radii
substantially larger than the original Solar nebula.
4) first planetesimals
Classic studies of planet formation by
accretion assume that planetesimals will congeal throughout an infant disk. Indeed,
many simulation studies of system evolution have started with a numerical set-up
in which protoplanets are packed in a series of concentric orbits separated by several
Hill radii (e.g., Raymond
et al. 2006, Ida & Lin 2010). Yet the ensemble of research discussed here
supports a different picture of accretion. It seems that planetesimals, the
essential precursors of protoplanets and planets, are born only in favored
locations: the planet traps and filters detailed above. System architectures
are thus erected on chemical and physical foundations established early in disk
evolution.
Dust, defined as particles smaller than one
millimeter, is evidently abundant throughout a newborn protoplanetary disk.
Dust grains can settle in the midplane and clump into pebbles, defined as
centimeter-sized particles. But further aggregation into planetesimals, defined
as particles measured in tens of meters, appears to be possible only in refuges
where gas drag and Type I migration can be overcome. Some models propose that
pebbles cluster in multitudes in these refuges, eventually thronging so densely
that they collapse into planetesimals 100 km in diameter – much larger than the
boulder-sized constituents invoked by older models. Objects this size can
readily accrete pebbles and grow even bigger.
As Alessandro Morbidelli & Sean Raymond (2016)
recently observed, “All this may be suggestive that the first planetesimals
formed only at select locations (near the disk’s inner edge or beyond the
snowline).” In this way, they say, the first generation of planetesimals “may serve
as a planetary system’s blueprint.”
5) dissipation & relaxation
A forming planetary system can be regarded as
a closed system engaged in feverish, time-limited activity. The processes of accretion,
migration, and orbital evolution are determined by factors originating within
the system: stellar irradiation, stellar outbursts, gas flow through the disk, planetesimal/disk
interactions, planetesimal/planetesimal interactions, planet/disk interactions,
and eventually planet/planet interactions. All these factors are sensitive to
the passage of time. For example, as gas flow slackens, migration slows, and as
gas dissipates, opportunities close down for young planets to accrete H/He
atmospheres. At the same time, the radial location of planet traps moves inward,
while their ability to trap solids diminishes.
The emergence of a gas giant planet anywhere
in the disk has far-reaching consequences. A gas giant clears gas and solids
from its orbit, cutting off the flow of volatiles and the migration of planetesimals
and protoplanets originating at wider radii. This development starves the inner
disk of replenishment for the gases, dust, and pebbles that it continually loses
to stellar accretion. The inner disk therefore drains rapidly into the star,
creating a central hole. Its removal exposes the outer disk to photoevaporation
by stellar flux. The process of dissipation is quite fast, measured in
thousands rather than millions of years.
Although not all disks give birth to giants,
most of them evidently evaporate on a similar schedule in a similar sequence: drainage
of the inner disk followed by photoevaporation of the outer disk. This might be
a universal process that is simply accelerated by the assembly of one or more
giant planets.
Disk dispersion is a watershed moment in
system evolution. Not only is gas accretion quenched for young planets; the
dissipation of H/He removes the mechanism by which orbital eccentricities are
damped, exposing planets to mutual perturbations. These interactions can happen
even in the presence of nebular gas, but their likelihood is significantly
amplified at this evolutionary phase. Orbit crossings, planetary collisions,
and planetary ejections are all possible, potentially leading to major
revisions in the “blueprint” created by primordial disk structures. Planet/planet
interactions after gas removal probably cause the orbital eccentricities
observed in multiplanet systems, while orbital sculpting and dynamical upsets
are especially likely in systems with gas giants outside 1 AU (Matsumura et al.
2013, Mustill et al. 2016).
External influences on system evolution are
also possible. Most stars shining today, including our Sun, formed in clusters
containing thousands of protostars packed within an area just a few parsecs in
radius. Under these conditions, stellar flybys are inevitable. Close encounters
in stellar nurseries can strip the outer disk while a star is accreting gases,
or scatter outer planets at any point during a system’s lifetime. Nevertheless,
studies disagree on the frequency of flybys (e.g., Malmberg et al. 2011, Li
& Adams 2015), so it is unclear how common or rare an extreme disruption
might be. In any case, given the short lifetimes of protoplanetary disks, flybys
are more likely to happen after the gas dissipates, when orbital architectures are
beginning to stabilize.
6) where do baby planets come from?
At the beginning of this series I asked the simple
question shown above. According to the body of research discussed here, the
simple answer seems to be: from farther
out in the protoplanetary disk. Baby planets originate in cooler zones
where the preservation and accretion of solids is supported at primordial times.
They achieve their mature compositions and orbits by interacting with local
solids and gases, typically engaging in orbital migration over some distance
and likely incorporating pebbles and planetesimals that formed in planet traps.
The disk models discussed here are unfriendly
to the accretion of protoplanets on short-period orbits during the earliest
phases of evolution. Instead, rocky objects assemble outside the silicate
condensation front and icy objects assemble in the snow region. As these fronts
move inward, their clutches of planetesimals and protoplanets are conveyed into
the inner system, contributing to the composition of planets with various
ratios of solids and volatiles. We can expect warm, low-mass planets that are
purely rocky, like the four terrestrial planets in the Solar System; part rock
and part ice, like the larger satellites of Jupiter and Saturn; or composed
primarily of rock and ice but supporting H/He envelopes, like Uranus and
Neptune. Gas giant planets, which seem constrained to form on still wider
orbits (in the snow region, near the heat transition, or even beyond the zone
where carbon monoxide freezes) often have opportunities to intrude on these
families of low-mass planets, with the odds of inner-system derangement
increasing along with the mass of the intruders.
Given these widely endorsed views, it surprises me that theoretical
models of planet composition in Kepler systems still limit their focus to rocky
objects with H/He envelopes but no water or other volatile content (e.g., Lopez
& Fortney 2014, Rogers 2015, Dorn et al. 2015, Wolfgang et al. 2016). I
look forward to studies that widen their scope to accommodate more diverse
internal structures.
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