Introduction — Nebular hypothesis
The nebular hypothesis, originally proposed by Immanuel Kant and later reformulated by Pierre Laplace, is the foundational cosmogonic model in which the Solar System originates from a rotating disk of gas and dust surrounding a young Sun. Its modern incarnation, the solar nebular disk model (SNDM), retains the core idea of a flattened, rotating protoplanetary disk and accounts naturally for the planets’ largely circular, coplanar orbits and their common direction of motion aligned with solar rotation, while many original mechanistic details have been superseded by subsequent theory.
Star and planet formation are linked processes that begin within giant molecular clouds dominated by molecular hydrogen. Gravitational instability fragments these clouds into denser, rotating clumps that collapse to form protostars; conservation of angular momentum produces a circumstellar accretion disk (a proplyd) as an inevitable byproduct, so planetary systems are a common consequence of star formation. The hierarchical observational sequence proceeds from diffuse interstellar medium to molecular cloud, Bok globule or dark nebula, young stellar object and protostar, and thence to pre–main-sequence stars such as T Tauri or Herbig Ae/Be objects; jets from these young stars generate Herbig–Haro objects where outflows interact with ambient gas.
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Timescales are constrained by both theory and observation: a Sun-like star assembles on the order of 10^6 years, its protoplanetary disk evolves and produces planetary bodies over 10^7–10^8 years, and the final accretional assembly of terrestrial planets by giant impacts may extend to ~10^8–10^9 years. The disk initially is hot and viscously accreting; during the T Tauri phase it cools sufficiently for refractory and volatile solid grains to condense. These grains coagulate by sticking and collisional growth into meter- to kilometer-scale planetesimals, the fundamental building blocks of planets.
Where local solids and dynamical conditions permit, runaway accretion can rapidly produce lunar- to Mars-sized planetary embryos on timescales of ~10^5–3×10^5 years. In the inner disk these embryos subsequently experience a protracted, stochastic phase of mutual collisions and mergers that yields a small number of terrestrial planets over tens to hundreds of millions of years. Beyond the frost line, lower temperatures permit ices to augment solid mass, enabling the formation of more massive cores—typically several times larger than inner-disk embryos and often reaching ~5–10 Earth masses—which can begin to accrete substantial hydrogen–helium envelopes.
Gas accretion onto sufficiently massive cores proceeds slowly at first over several million years, but when a protoplanet attains roughly 30 Earth masses a runaway gas-accretion regime is triggered and the envelope grows very rapidly. Classical gas giants are thought to acquire most of their massive gaseous envelopes during this brief rapid phase (potentially as short as ~10^4 years), with growth terminating when the disk gas is dissipated. If cores form too late, after much of the nebular gas has been lost, they remain as ice-rich, lower-mass planets—an explanatory framework for Uranus- and Neptune-like ice giants as “failed” gas giants.
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Planetary migration—radial orbital evolution driven by disk–planet interactions or later dynamical processes—plays a central role in shaping final system architectures by relocating planets from their formation sites and by mediating interactions that depend sensitively on the timing of disk dispersal. The contemporary synthesis of planet formation therefore weaves together key theoretical concepts: accretional growth, the initial mass function of stars, Jeans instability for gravitational collapse, Kelvin–Helmholtz contraction and cooling during protostellar/protoplanetary evolution, the nebular/disk origin of planetary systems, and orbital migration. Together these elements form the conceptual framework of the modern nebular model.
History
The nebular hypothesis originated in the eighteenth century with early suggestions that rotating clouds of gas could condense to form planetary systems. Emanuel Swedenborg articulated elements of this idea in 1734, and Immanuel Kant elaborated it in 1755 by proposing that extended gaseous nebulae undergo slow rotation, gravitational collapse and flattening, producing condensed bodies that become stars and planets. Pierre‑Simon Laplace independently advanced a more detailed mechanism in 1796: a hot protosolar atmosphere contracts and spins up, shedding successive gaseous rings that cool and condense into planets. Laplace’s ring‑ejection formulation dominated nineteenth‑century cosmology as the canonical account of Solar System origin.
Despite its influence, the Laplacian model confronted a fundamental dynamical difficulty: it could not account for the observed distribution of angular momentum in the Solar System, in which the planets possess the vast majority of the system’s angular momentum while the Sun contains most of the mass. This mismatch led many astronomers to reject the ring hypothesis by the early twentieth century. Later historical reassessment has also corrected a persistent error in the secondary literature: a commonly cited nineteenth‑century technical critique—sometimes attributed to James Clerk Maxwell—that ring differential rotation would preclude condensation does not accurately reflect Maxwell’s published arguments, an attribution distorted by later popularizations.
Contemporary nineteenth‑century critiques included more publicly pitched objections. For example, Sir David Brewster argued against the plausibility of ring ejection on physical and empirical grounds, contending that a moon formed from such a ring ought to retain volatiles and an atmosphere, and noting philosophical resistance (exemplified by Newton) to spontaneous system formation without design. In the twentieth century a succession of alternative models sought to resolve Laplace’s problems: Chamberlin and Moulton’s planetesimal theory (1901), Jeans’s tidal hypothesis (1917), Schmidt’s accretion ideas (1944), McCrea’s protoplanet concept (1960), capture theories such as Woolfson’s, and Prentice’s 1978 revival of Laplacian motifs. None of these proposals produced a fully satisfactory, quantitatively robust account; many remained largely descriptive.
The modern consensus—the solar nebular disk model (SNDM)—was shaped decisively by Victor Safronov’s 1969 monograph, which systematically identified the principal problems of planetary formation and proposed mechanisms that redirected subsequent research. Work by later theorists, notably George Wetherill’s demonstration of runaway accretion as an efficient route to rapid planetary embryo growth, transformed the SNDM into a comprehensive theoretical framework applicable beyond the Solar System. Empirical advances have reinforced this paradigm: large exoplanet surveys have revealed thousands of planetary systems (6,032 confirmed extrasolar planets within the Galaxy as of 29 July 2025), supporting the view that disk‑mediated planet formation is a common astrophysical process.
Achievements (Nebular Hypothesis)
Observations and theory conjointly establish that star formation almost invariably produces circumstellar accretion disks around young stellar objects—protostars and T Tauri stars—so that disks are intrinsic to early stellar evolution. By an age of order 10^6 years essentially all young stars retain gaseous and dusty disks, consistent with both direct imaging and theoretical expectations. High-resolution studies reveal extremely rapid dust processing within these disks: submicron grains coagulate into macroscopic solids on very short (∼10^3 yr) timescales, yielding centimeter-sized particles in brief intervals. Growth continues hierarchically from these pebbles to kilometre-scale planetesimals, and onward to larger bodies; the stage that converts ∼1 km planetesimals into 10^3 km objects is well characterized and requires locally high solid densities to proceed efficiently. Early assembly is dominated by runaway accretion—where the largest objects grow fastest—followed by an oligarchic phase in which a population of dominant embryos accretes more slowly. The resulting planetary embryos exhibit characteristic sizes that vary systematically with orbital radius, reflecting the disk’s local physical conditions and surface density of solids. Numerical simulations of embryo–embryo interactions and collisions in the inner disk reproduce the formation of a small number of Earth-sized bodies, providing a quantitatively supported pathway for terrestrial-planet formation. Taken together—ubiquitous disks, rapid grain growth, and a well-constrained accretionary sequence from dust to planet-sized bodies—these results render the origin of terrestrial planets largely understood within current theoretical and observational frameworks.
Current issues
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A central dynamical problem for protostellar accretion disks is removal of angular momentum: gas that ultimately accretes onto the central star must shed most of its specific angular momentum. Early proposals invoked angular-momentum loss in a stellar wind during the T Tauri phase, with viscous stresses in the disk transporting the corresponding angular momentum outward. In practice, disk evolution models therefore rely on an effective viscosity produced by macroscopic turbulence, but the microphysical origin of that turbulence remains inadequately constrained. Magnetic braking—transfer of stellar spin into the surrounding disk via magnetic coupling—provides a competing or complementary pathway for extracting angular momentum.
The dispersal of the gaseous component of protoplanetary disks is controlled principally by two processes: viscous diffusion, which redistributes gas radially and feeds accretion onto the star, and photoevaporation, in which radiative heating drives thermal winds that remove gas. The interplay of these mechanisms largely determines disk lifetimes; observational surveys indicate characteristic gas-disk lifetimes of order <10^7 years, a timescale that historically posed a challenge for planet-formation models.
Two major formation problems remain unresolved or only recently alleviated. The formation of planetesimals—the transition from centimeter- to kilometer-scale solids—constitutes a key gap in theory: there is no universally accepted, empirically verified mechanism that explains how metric- to decimeter-sized particles overcome rapid radial drift and destructive collisions to grow into kilometer-scale bodies. This unresolved step likely controls which systems retain substantial solid reservoirs and which remain largely devoid of large bodies or debris belts. Separately, the apparent timing conflict between giant-planet core growth and disk lifetimes has been relaxed by theoretical advances; contemporary models can assemble Jupiter-mass planets in of order 4 × 10^6 years or less, bringing giant-planet formation into concordance with observed gas-disk lifetimes.
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Interactions between forming planets and their natal disks introduce an additional dynamical difficulty: differential torques can drive rapid inward migration of massive cores and nascent giants, potentially delivering them into the inner system before they accrete substantial envelopes. Recent work that treats the concurrent evolution of the disk—its thermodynamics, mass redistribution, and dissipative processes—shows that migration can be slowed, reversed, or halted in many circumstances, mitigating the most extreme inward drift predicted by simpler models. Observed multiple and dynamically interacting systems (for example AS 205) underscore the complexity of these processes and the need to couple angular-momentum transport, disk dispersal, planetesimal formation, and planet–disk interactions within a unified evolutionary framework.
Protostellar formation proceeds within giant molecular clouds—cold, primarily molecular‑hydrogen reservoirs whose typical masses (~3×10^5 M☉) and diameters (~20 pc) supply the material for clustered star formation. Observations of regions such as the Trifid Nebula illustrate this environment: visible‑light images show dusty obscuration, whereas infrared imaging penetrates the dust to reveal embedded star‑forming activity otherwise hidden at shorter wavelengths.
Over millions of years such clouds gravitationally collapse and fragment into numerous dense cores (protostellar or protosolar nebulae) with masses from sub‑solar to several solar masses. Individual cores are compact (diameters ~0.01–0.1 pc, i.e., ~2,000–20,000 AU) and dense (particle number densities ~10^4–10^5 cm−3), conditions that both facilitate collapse and shield nascent protostars from external radiation. The initial collapse of a solar‑mass core proceeds on a characteristic timescale of order 10^5 years, producing a small, hot hydrostatic core that contains only a minor fraction of the original mass and seeds the future star.
Conservation of angular momentum during infall spins up the envelope and prevents direct deposition of all material onto the central object; instead matter settles into a rotationally supported circumstellar disk. At the earliest (Class 0) stage the protostar+disk system remains enshrouded by an optically thick envelope so opaque that even millimeter‑wave emission can be heavily attenuated, and these objects are therefore detected primarily as bright millimeter–submillimeter condensations. Their luminosity is dominated by release of gravitational energy, and a solar‑mass Class 0 protostar can attain luminosities up to ~100 L☉. Collapse commonly launches bipolar jets and molecular outflows along the rotation axis; Herbig–Haro objects and well‑studied molecular flows (e.g., the infrared outflow associated with HH 46/47) are observable manifestations of this activity.
As envelope infall diminishes, the system becomes detectable first in the far‑infrared and later at optical wavelengths, marking the transition to the Class I young stellar object. Deuterium fusion begins at this stage and, if the central mass exceeds ≈80 Jupiter masses (~0.08 M☉), sustained hydrogen fusion will later commence; objects below this threshold become brown dwarfs. This evolutionary step occurs roughly 10^5 years after the onset of collapse, by which time the combined mass of the disk and residual envelope is typically ≤10–20% of the central object’s mass.
After about 10^6 years the envelope has largely been accreted or dispersed and the system appears as a classical T Tauri star surrounded by a protoplanetary disk whose mass is typically ~1–3% of the stellar mass. Accretion rates at this stage are of order 10−7 to 10−9 M☉ yr−1; magnetospheric accretion produces strong emission lines, photometric variability, magnetic activity, and continued bipolar jets. Over the next ~10^7 years the disk is depleted by a combination of accretion onto the star, planet formation, ejection in jets, and photoevaporation by stellar and external UV radiation. The object then evolves into a weak‑line T Tauri star and, over hundreds of millions of years, settles onto the main sequence as a Sun‑like star.
Protoplanetary disks
Protoplanetary (debris) disks are an almost ubiquitous feature of early stellar evolution, typically arising concurrently with the protostar and observable around a large fraction of stars in young clusters. High-resolution and reprocessed archival observations continue to reveal previously hidden circumstellar material (for example, debris disks recovered around HD 141943 and HD 191089 in HST data), illustrating both the observational diversity of disks and the value of improved imaging techniques. During the earliest (Class 0) phase the dense protostellar envelope often obscures the innermost disk, but theoretical expectations and indirect evidence indicate the presence of massive, hot accretion disks that feed the growing star.
Thermal structure and composition vary strongly with radius. Viscous dissipation of turbulence and energy released by infalling nebular gas heat the inner disk: model estimates place temperatures above ~400 K within ∼5 AU and near 1,000 K inside ∼1 AU. These high temperatures vaporize volatile species (water, organics and some silicates), leaving refractory materials such as iron concentrated toward the center, while ices and more volatile compounds survive only in the outer disk. As the system evolves toward the classical T Tauri stage (typical disk lifetimes ≲10 Myr) the disk thins and cools, permitting condensation of less-volatile minerals into submicron-to-micron grains, including crystalline silicates.
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Transport processes — radial and vertical mixing together with turbulent viscosity — play a central role in disk evolution. Turbulent angular-momentum transport permits mass to flow inward onto the protostar while a small fraction of gas moves outward carrying away angular momentum, allowing the disk to grow in radius; if the initial nebular angular momentum is high, observed disks can extend to ≳1,000 AU. Mixing also redistributes newly condensed crystalline grains and primordial, volatile-rich grains, providing a natural pathway for the presence of interstellar components and refractory inclusions in primitive meteorites and cometary material.
Solid growth proceeds from submicron grains to aggregates and then to planetesimals, but the pathway beyond centimeter scales is problematic. In dense disk regions submicron particles coagulate into aggregates up to several centimeters — a process traced in infrared spectra of young disks — and continued growth can produce kilometer-scale planetesimals, the basic building blocks of planets. Simple collisional sticking becomes inefficient as sizes increase, so leading hypotheses for planetesimal formation invoke collective and gravitational processes. One idea is that solids settle to the midplane forming an exceedingly thin (~<100 km) dense layer that fragments by gravitational instability into planetesimals; in practice, shear-driven turbulence often prevents such extreme settling except where solids are locally enhanced. The streaming instability offers an alternative: aerodynamic feedback concentrates particles by reducing the headwind from sub-Keplerian gas, allowing clumps to grow into dense filaments that gravitationally collapse into planetesimals comparable to large asteroids.
Formation of gas giants by direct fragmentation of the gaseous disk (disk-scale gravitational instability) is possible but limited. Collapse of gaseous clumps can form massive planets or brown dwarfs on dynamical timescales (~10^3 years) only if the disk mass is substantial (≳0.3 M☉); by contrast, typical observed disk masses are an order of magnitude or more lower (≈0.01–0.03 M☉), rendering this route relatively rare but potentially important for brown-dwarf formation.
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Disk dispersal is the outcome of multiple, concurrent processes: inner gas is accreted onto the star or expelled in bipolar jets; photoevaporation by stellar UV (or by nearby massive stars) removes outer gas; forming planets accrete or scatter gas; and small dust can be blown out by radiation pressure. The end states range from full planetary systems to remnant dust disks or systems with little observable circumstellar material if planetesimal formation is inefficient. Surviving planetesimals become the asteroids and comets of mature systems; meteorites sample both primitive, undifferentiated bodies and fragments of thermally processed, differentiated planetesimals, and young systems can also capture interstellar objects. Well-studied examples spanning these phenomena include the disks imaged around HD 141943 and HD 191089, the complex disk and gas streams inferred for HD 142527, and the active planetesimal/exocomet environment of Beta Pictoris, which together illustrate the range of morphologies and evolutionary pathways in protoplanetary disks.
Rocky planets form within the inner region of a protoplanetary disk inside the stellar frost line (roughly the inner 3–4 AU for a Sun-like star), where temperatures prevent condensation of water ice and other volatiles. In this hot inner zone solid growth proceeds by coagulation of refractory grains, producing rocky planetesimals that seed subsequent accretion.
Once planetesimals reach sizes of order 1 km, runaway accretion ensues: larger bodies grow ever more rapidly because accretion rates scale strongly with size (approximately ∝ R^4, or ≈ M^{4/3}). This phase is short-lived (∼10^4–10^5 yr) and typically terminates when the largest objects reach diameters near 1,000 km; their strong gravitational perturbations stir and heat the remaining small bodies, reducing subsequent accretion efficiency.
Growth then transitions to an oligarchic regime characterized by a few hundred dominant embryos (oligarchs) that slowly sweep up the remaining planetesimals while preventing rival bodies from substantial growth. In this stage an oligarch’s accretion rate is set mainly by its geometric cross section (∝ R^2), so the specific accretion rate falls with mass (∝ M^{−1/3}), allowing smaller oligarchs to partially catch up. Oligarchs remain spaced by roughly 10 times their Hill radius, Hr = a(1 − e)(M/3Ms)^{1/3}, and they occupy low-eccentricity, low-inclination orbits within cleared gaps separated by rings of leftover planetesimals.
An oligarch’s final mass (the isolation mass) depends on radial distance and the local surface density of solids; in the inner disk this isolation mass reaches up to ~0.1 M⊕ (≈ one Mars mass). Oligarchic evolution therefore produces on the order of 100 Moon- to Mars-sized embryos, spaced at ≈10·Hr, over a timescale of a few 10^5 years. When the reservoir of small planetesimals is depleted and embryos become sufficiently massive, mutual perturbations drive the system into a chaotic merger or “giant-impact” stage.
The merger stage, lasting roughly 10–100 Myr, is dominated by close encounters, collisions and ejections: embryos clear remaining small bodies by scattering and then collide with one another until only a few large planets remain. Numerical experiments typically produce 2–5 surviving Earth-sized planets per system; many embryos merge to form these planets while a comparable number are lost to ejection. In the Solar System this picture explains several features: Earth and Venus plausibly accreted from ~10–20 embryos each (with a similar number ejected), some embryos originating in the asteroid belt may have delivered water to the growing Earth, and Mars and Mercury can be interpreted as embryos that avoided complete merger.
The final rocky planets generally settle into relatively long-lived orbits but remain dynamically “packed,” often near the threshold of instability. This near-limit packing is a natural outcome of the oligarchic and merger phases and helps account for the compact architectures observed in many planetary systems.
Giant planets
Observations of circumstellar debris, exemplified by the asymmetric dust ring around Fomalhaut, lend empirical support to the idea that massive planets sculpt protoplanetary disks: gaps and non‑axisymmetric features are readily explained by the gravitational perturbations of one or more giant companions. Within the nebular framework two principal formation pathways have been proposed. In the disk‑instability picture a sufficiently massive, cool disk fragments directly under its own gravity, whereas the core‑accretion (nucleated instability) model builds a solid nucleus that subsequently captures a massive gaseous envelope. Surveys up to 2011 favor core accretion as the dominant route, although either mechanism can, under extreme conditions, produce objects with brown‑dwarf masses.
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Core accretion operates in two broad phases. Solid material first coagulates into a ~10 M⊕ core through the familiar sequence of runaway growth followed by an oligarchic regime; beyond the frost line this process is aided by abundant ices, which raise the available solid mass by roughly a factor of four relative to purely rocky composition (ice:rock ≈ 4:1). Once a core attains of order 5–10 M⊕ it begins to bind a substantial gaseous envelope. An initial quasi‑static envelope growth can increase the combined mass to ∼30 M⊕ over a few million years, after which runaway gas accretion ensues and most of the planet’s final mass is accumulated on timescales of order 10^4 years. Accretion halts when the planet exhausts its local gas supply, opens a gap by tidal interaction, or the global disk disperses; if dispersal precedes runaway accretion, the object remains an ice giant rather than a gas giant.
Quantitative problems arise under a Minimum Mass Solar Nebula (MMSN): at Jupiter’s orbital distance (≈5 AU) the MMSN supplies only ≈1–2 M⊕ of solids within typical gas‑disk lifetimes (~10 Myr), too little to form a giant‑planet core rapidly by classical planetesimal accretion. Proposed resolutions include invoking more massive disks (order‑of‑magnitude increases), inward migration of embryos that permits them to sample and accrete more solids, and enhanced accretion efficiency caused by gas drag within growing envelopes. A distinct and highly effective channel is pebble accretion: centimeter–meter sized particles whose radial drift is moderated by aerodynamic drag can be captured very efficiently by embryos, accelerating core growth by orders of magnitude (potentially up to ~10^3 times faster than classical planetesimal accretion) and easing formation at large distances.
Forming Uranus‑ and Neptune‑class planets in situ at ~20–30 AU remains difficult under standard core‑accretion assumptions; plausible alternatives are formation closer in followed by scattering and outward migration, or very rapid pebble accretion at large orbital radii. During active gas accretion a circumplanetary subdisk typically develops, mediating the supply of solids and gas to the planet and providing the birthplace for regular satellite systems such as the Galilean moons.
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Giant planets also exert a profound dynamical influence on inner planetary systems. Their torques on the gas disk drive orbital migration, which can carry giants inward until halted at an inner truncation and produce the observed population of hot Jupiters. Gravitational perturbations from forming giants raise eccentricities and inclinations of planetesimals and embryos within several AU, altering terrestrial planet assembly: early giant formation can suppress inner accretion, whereas late formation tends to make embryo collisions more energetic, reducing the final number of planets while increasing their typical masses and compactness. In the immediate vicinity of a giant planet embryos are strongly excited onto high‑eccentricity orbits, leading to close encounters and frequent ejection; although giants alone are inefficient at clearing small bodies, the combined action of embryos and giants typically removes ≈99% of the initial planetesimal population, leaving a low‑mass remnant that can evolve into an asteroid‑belt analogue.
Exoplanets
Over the last two decades thousands of planets have been detected around other stars, and statistical extrapolation implies that billions more exist within the observable universe. This census reveals orbital architectures far richer than the Solar System’s, including very close-in giant planets, intermediate-mass rocky bodies, and tightly packed multi-planet systems with short orbital periods.
Observed classes cluster around a few common types. “Hot” and “warm” Jupiters are gas giants on close or moderately close orbits; super-Earths occupy the mass range above Earth but well below Neptune; and many systems comprise multiple short-period planets arranged in compact configurations. These empirical groupings reflect distinct formation and dynamical histories rather than a single pathway.
Several mechanisms can place planets on close-in or eccentric orbits. During the gaseous protoplanetary phase, planets exchange angular momentum with the disk and migrate inward: low-mass bodies primarily undergo Type I migration while gap-opening massive planets follow the slower Type II mode, a process capable of delivering giant planets from distant formation sites to warm or hot orbits. Alternatively, strong gravitational interactions among planets can scatter a body onto a highly eccentric orbit whose periastron lies near the star; tidal dissipation at each close approach then removes orbital energy and circularizes the orbit at small semimajor axis. Secular perturbations from a distant, inclined companion can also drive large eccentricity oscillations (Kozai–Lidov cycles), allowing subsequent tidal circularization to produce close-in giants. The common occurrence of significant eccentricities among Jupiter-sized exoplanets points to past planet–planet encounters in many systems, though eccentricity can also grow during resonant convergent migration.
The origin of hot Jupiters and their cores may therefore be ex situ (formation at larger radii followed by inward transport) or in situ (local accretion near present orbits). In situ formation requires a substantial reservoir of solids close to the star or the prior inward delivery of cores that then accrete gas; cores themselves might have formed locally or migrated inward before runaway gas capture. For lower-mass planets, multiple pathways also operate. Super-Earths can assemble in place if the inner disk is sufficiently massive, or they can result from inward migration and coalescence of embryos or from the radial drift and concentration of small solids from the outer disk. Because these bodies are low mass, their migration is expected to be dominated by Type I torques; the presence of resonant orbital chains in some systems provides evidence for past migration, whereas widely spaced, non-resonant systems are more consistent with a phase of post-disk dynamical instability that broke initial resonances.
Giant planets exert a strong sculpting influence on inner-system architecture. An early-forming giant like Jupiter can block the inward flux of solids and embryos, inhibiting the formation or emplacement of a compact inner system—an effect that may help explain the paucity of close-in super-Earths in the Solar System.
The acquisition of primordial gas envelopes by super-Earths depends sensitively on the timing of giant impacts relative to disk dispersal. If major mergers occur after the gas has dissipated, rocky, atmosphere-poor planets result; mergers during a disk transition stage can leave a modest (percent-level) envelope; if assembly completes while the disk remains dense, runaway gas accretion can produce a gas giant. The onset of collision-dominated growth is governed by the reduction of gas-driven dynamical friction: as disk damping weakens, orbits cross and giant impacts commence, a transition that occurs earlier in higher-metallicity disks because solids and embryos grow more rapidly. Finally, gas accretion onto cores can be self-limiting if envelope structure permits substantial gas flow through rather than retention, delaying runaway accretion until core masses approach roughly 15 Earth masses, beyond which rapid gas capture becomes likely.
Together, these processes—disk-mediated migration, dynamical scattering and secular excitation, timing of collisions and gas dissipation, and the modulating role of giant planets and disk metallicity—produce the observed diversity of exoplanet masses, compositions and orbital configurations.
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In planet-forming systems the term “accretion disk” is used in two distinct senses that should be distinguished. One sense denotes a gaseous disk that channels mass inward onto a central contracting object (for example a T Tauri–type protostar): viscous stresses and associated angular‑momentum transport produce a net radial flux of gas from larger to smaller radii, delivering material onto the stellar surface. The other sense denotes the local environment within a protoplanetary disk where solid particles grow by collisional and gravitational aggregation: cooled dust and ice grains collide, stick, and agglomerate into larger aggregates, then planetesimals, and ultimately planetary bodies through a succession of low‑ and high‑energy impacts moderated by aerodynamic forces and gravitational focusing.
These two classes of accretion operate under different physical regimes. Gas accretion is governed by fluid dynamics and angular‑momentum transport (viscosity, turbulence, magnetic effects), whereas planetary accretion among solids is controlled by particle aerodynamics, collision physics, coagulation, and N‑body gravity. The same distinction applies at smaller scale: forming giant planets are surrounded by circumplanetary gaseous disks that channel hydrogen–helium onto the protoplanet, while solid material within those disks undergoes local coagulation to assemble the regular satellite systems.
For conceptual clarity and for accurate modeling it is therefore important to treat inward gas‑feeding processes and solid‑body coagulation as separate phenomena. Although they coexist spatially within disk environments, they involve different materials, mechanisms, timescales, and dynamical outcomes—feeding central masses in the first case versus building orbiting solids in the second.