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Gas and ice giants

Growth of massive cores beyond the frost line, attracting thick hydrogen–helium envelopes

1. Beyond the frost line

In protoplanetary disks, in the region located beyond a certain distance from the star – often called the frost line (snow line) – water and other volatile substances can freeze into ice grains. This has a significant impact on planet formation:

  1. Ice-enriched solid particles: Lower temperatures allow water, ammonia, methane, and other volatile substances to condense onto dust grains, increasing the total mass of solid material.
  2. Larger solid particle cores: This increase in mass helps planetary embryos rapidly accumulate material and reach a critical mass to attract nebular gas.

Because of this, planets forming in the outer part of the disk can acquire thick hydrogen–helium envelopes and become gas giants (like Jupiter or Saturn) or ice giants (Uranus and Neptune). While terrestrial planets in the hot inner disk remain relatively low in mass and mostly rocky, these outer disk planets can reach tens or hundreds of Earth masses, significantly influencing the overall planetary architecture of the system.

2. Core accretion model

2.1 Basic premise

The widely accepted core accretion model states:

  1. Solid core growth: The planetary embryo (initially an ice-enriched protoplanetary body) accretes local solids until reaching ~5–10 MEarth.
  2. Gas accretion: When the core becomes massive enough, it rapidly gravitationally attracts hydrogen–helium from the disk, initiating runaway envelope accretion.
  3. Runaway growth: This forms Jupiter-type gas giants or intermediate-sized “ice giants” if disk conditions are less favorable for envelope accretion or the disk disperses earlier.

This model reliably explains the existence of massive H/He envelopes around jovian planets and more modest envelopes in “ice giants,” which may have formed later, accreted gas more slowly, or lost part of their envelope due to stellar or disk processes.

2.2 Disk lifetime and rapid formation

Gas giants must form before the disk gas disperses (within ~3–10 million years). If the core grows too slowly, the protoplanet will not have time to accrete much hydrogen–helium. Studies in young star clusters show that disks dissipate quite rapidly, correspondingly supporting that giant planet formation must occur quickly enough to utilize the short-lived gas reservoir [1], [2].

2.3 Envelope contraction and cooling

Once the core surpasses the critical mass, an initially shallow atmospheric layer transitions into a runaway gas accretion phase. As the envelope grows, gravitational energy is radiated away, allowing the envelope to contract and attract even more gas. This positive feedback can form final planets of ~tens to hundreds of Earth masses, depending on local disk density, time, and factors such as Type II migration or gap formation in the disk.

3. Frost lines and the importance of icy solids

3.1 Volatile compounds and increased solid particle mass

In the outer disk, where the temperature drops below ~170 K (for water, although the exact limit depends on disk parameters), water vapor condenses, increasing the surface density of solid particles by 2–4 times. Other ices (CO, CO2, NH3) also condense at even lower temperatures further from the star, so the amount of solid material becomes even greater. This abundance of ice-enriched planetesimals leads to faster-growing cores, which is the main premise for the formation of gas and ice giants [3], [4].

3.2 Why do some become gas giants while others become ice giants?

  • Gas giants (e.g., Jupiter, Saturn): Their cores form quickly enough (>10 Earth masses) to capture a massive hydrogen–helium layer from the disk.
  • Ice giants (e.g., Uranus, Neptune): May have formed later, accreting more slowly or experiencing greater disk dispersion, resulting in a smaller gas envelope, with much of their mass composed of water/ammonia/methane ices.

Thus, whether a planet becomes a "Jovian giant" or a "Neptunian ice giant" depends on the density of solid particles, core growth rate, and external environment (e.g., photoevaporation from nearby massive stars).

4. Growth of massive cores

4.1 Planetesimal accretion

Based on the rigid core accretion model, icy planetesimals (km-sized and larger) form through collisions or streaming instability. When a protoplanet reaches ~1000 km or larger, it enhances gravitational collisions with remaining planetesimals:

  1. Oligarchic growth: Several large protoplanets dominate the region, "sweeping up" smaller body populations.
  2. Fragmentation reduction: Lower collision speeds (due to partial gas damping) promote accretion rather than disruption.
  3. Timescales: The core must reach ~5–10 MEarth within a few million years to take advantage of the disk gas [5], [6].

4.2 "Pebble" accretion

Another mechanism is "pebble" accretion:

  • Pebbles (mm–cm) drift through the disk.
  • A sufficiently massive protoplanetary core can gravitationally "capture" those pebbles, growing very rapidly.
  • This accelerates the transition to a super-Earth or giant core, which is crucial to initiate envelope accretion.

When the core reaches the critical mass, runaway gas accretion begins, resulting in a gas giant or ice giant depending on the final envelope mass and disk conditions.

5. Envelope accretion and gas-dominated planets

5.1 Uncontrolled envelope growth

Once the core exceeds the critical mass, the pro-giant planet initially has a thin atmosphere that transitions into an uncontrolled gas accretion phase. As the envelope expands, gravitational energy is radiated away, allowing even more nebular gas to be attracted. A key limiting factor is often the disk's ability to supply and replenish gas or the planet's ability to cool and attract its envelope. Models show that if a ~10 MEarth core forms, the envelope mass can grow to tens or hundreds of Earth masses if the disk remains [7], [8].

5.2 Gap formation and Type II migration

A sufficiently massive planet can clear a gap in the disk through tidal torques exceeding local disk pressure forces. This alters gas supply and leads to Type II migration, where the planet's orbital evolution depends on the disk's viscosity scale. Some giants may migrate inward (forming "hot Jupiters") if the disk does not dissipate quickly enough, while others remain in or beyond their formation zone if disk conditions suppress migration or if multiple giants lock in resonances.

5.3 Various final outcomes of gas giants

  • Jupiter-like: Very massive, large envelope (~300 Earth masses), ~10–20 Earth mass core.
  • Saturn-like: Intermediate envelope size (~90 Earth masses), yet clear hydrogen–helium dominance.
  • Sub-Jovians: Lower total mass or incomplete uncontrolled growth.
  • Brown dwarfs: At about ~13 Jupiter masses, the boundary arises between giant planets and substellar brown dwarfs, although formation mechanisms may differ.

6. Ice giants: Uranus and Neptune

6.1 Formation in the outer disk

Ice giants like Uranus and Neptune have a total mass of about 10–20 Earth masses, with ~1–3 MEarth in the core and only a few Earth masses in the hydrogen/helium envelope. They are thought to have formed beyond 15–20 AU, where disk density is lower and accretion rates slow due to greater distance. The causes of their formation differ from Jupiter/Saturn:

  • Late formation: The core reached critical mass quite late, while the disk was already dispersing, resulting in less gas accreted.
  • Faster disk dissipation: Less time or external radiation reduced gas reserves.
  • Orbital migration: They may have formed somewhat closer or farther and been pushed to current orbits due to interactions with other giants.

6.2 Composition and internal structure

Ice giants contain abundant water/ammonia/methane ices — volatile compounds that condensed in the cold outer zone. Their higher density compared to pure H/He giants indicates more "heavy elements." The internal structure may be layered: a rocky/metallic core, a hydrogen mantle with dissolved ammonia/methane, and a relatively thin H–He layer on top.

6.3 Exoplanetary analogs

Many exoplanets called "mini-Neptunes" have masses intermediate between super-Earths (~2–10 MEarth) and Saturn. This indicates that partial or incomplete envelope accretion is quite common once at least a medium-sized core forms — a dynamic similar to the formation of an "ice giant" around many stars.

7. Observational tests and theoretical considerations

7.1 Observing forming giants in disks

ALMA-detected ring/gap patterns may be carved by giant planet cores. Some direct imaging instruments (e.g., SPHERE/GPI) attempt to detect young giant objects still embedded in the disk. Such detections confirm the pulls and mass accumulation predicted by core accretion theory.

7.2 Composition clues from atmospheric spectra

Spectra of exoplanet giants (from transit or direct observations) reveal atmospheric “metallicity,” indicating how many heavy elements are present. Observing Saturn's and Jupiter's atmospheres also shows traces of disk chemistry when they formed, e.g., C/O ratio or noble gas amounts. Differences may indicate planetesimal accretion or dynamic migration paths.

7.3 Effects of migration and system architecture

Exoplanet surveys show many systems with hot Jupiters or multiple jovian planets close to the star. This indicates that giant planet formation and disk or planet-planet interactions can strongly shift orbits. Our Solar System's outer gas/ice giants shaped the final arrangement by scattering comets and smaller bodies, possibly helping protect Earth from greater migration threats (e.g., Jupiter or Saturn moving inward).

8. Cosmological consequences and diversity

8.1 Influence of stellar metallicity

Stars with higher metallicity (a higher fraction of heavy elements) generally more often have giant planets. Studies show a strong correlation between stellar iron abundance and the likelihood of giant planets. This is likely related to a greater amount of dust in the disk, which accelerates core growth. Low-metallicity disks often form fewer or smaller giants, or perhaps more rocky/“ocean” worlds.

8.2 The “brown dwarf desert”?

When gas accretion reaches about 13 Jupiter masses, the boundary between giant planets and substellar brown dwarfs becomes unclear. Observations show a “brown dwarf desert” near Sun-like stars (brown dwarfs are rarely found at close distances), possibly because bodies of such mass follow a different formation mechanism, and disk fragmentation rarely produces stable orbits in that mass range.

8.3 Low-mass stars (M dwarfs)

M dwarfs (lower mass stars) generally have lower mass disks. It's easier to form mini-Neptunes or super-Earths there than Jupiter-sized planets, although exceptions exist. The relationship between disk mass and stellar mass explains why Neptunes or rocky super-Earths are more commonly found around smaller stars.

9. Conclusion

Gas and ice giants are among the most massive outcomes of planetary formation, arising beyond the snow line in protoplanetary disks. Their powerful cores, rapidly formed from ice-enriched planetesimals, attract thick hydrogen–helium envelopes while the disk is gas-rich. The final results – jovian giants with huge envelopes, Saturn analogs adorned with rings, or smaller “ice giants” – depend on disk properties, formation pace, and migration history. Observations of exoplanet giants and gaps in young dusty disks show this process is widespread, driving orbital and compositional diversity in giant planets.

According to the core accretion model, the path appears nuanced: an ice-enriched body crosses several Earth masses, triggers runaway gas accretion, and becomes a massive H/He reservoir, largely shaping the entire planetary system's layout – dispersing or organizing smaller bodies, creating the main dynamical context. As we continue to observe ALMA ring structures, giant atmospheres' spectral data, and exoplanet statistics, our understanding of how cold protoplanetary disk zones grow the largest members of planetary families deepens.

Links and further reading

  1. Pollack, J. B., et al. (1996). “Formation of the Giant Planets by Concurrent Accretion of Solids and Gas.” Icarus, 124, 62–85.
  2. Safronov, V. S. (1972). Evolution of the Protoplanetary Cloud and Formation of the Earth and Planets. NASA TT F-677.
  3. Lambrechts, M., & Johansen, A. (2012). “Rapid growth of gas-giant cores by pebble accretion.” Astronomy & Astrophysics, 544, A32.
  4. Helled, R., et al. (2014). “Giant planet formation, evolution, and internal structure.” Protostars and Planets VI, University of Arizona Press, 643–665.
  5. Stevenson, D. J. (1982). “Formation of the giant planets.” Annual Review of Earth and Planetary Sciences, 10, 257–295.
  6. Mordasini, C., et al. (2012). “Characterization of exoplanets from their formation. I. Models of combined planet formation and evolution.” Astronomy & Astrophysics, 541, A97.
  7. Bitsch, B., Lambrechts, M., & Johansen, A. (2015). “The growth of planets by pebble accretion in evolving protoplanetary discs.” Astronomy & Astrophysics, 582, A112.
  8. D’Angelo, G., et al. (2011). “Extrasolar planet formation.” Exoplanets, University of Arizona Press, 319–346.
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