Sunday, November 27, 2016

Accretion with Migration in Radially Structured Disks

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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