Inastrophysics,accretion is the accumulation of particles into a massive object bygravitationally attracting more matter, typicallygaseous matter, into anaccretion disk.[1][2] Mostastronomical objects, such asgalaxies,stars, andplanets, are formed by accretion processes.
The accretion model that Earth and the otherterrestrial planets formed from meteoric material was proposed in 1944 byOtto Schmidt, followed by theprotoplanet theory ofWilliam McCrea (1960) and finally thecapture theory ofMichael Woolfson.[3] In 1978,Andrew Prentice resurrected the initial Laplacian ideas about planet formation and developed themodern Laplacian theory.[3] None of these models proved completely successful, and many of the proposed theories were descriptive.
The 1944 accretion model by Otto Schmidt was further developed in a quantitative way in 1969 byViktor Safronov.[4] He calculated, in detail, the different stages of terrestrial planet formation.[5][6] Since then, the model has been further developed using intensive numerical simulations to studyplanetesimal accumulation. It is now accepted that stars form by the gravitational collapse ofinterstellar gas. Prior to collapse, this gas is mostly in the form of molecular clouds, such as theOrion Nebula. As the cloud collapses, losing potential energy, it heats up, gaining kinetic energy, and the conservation ofangular momentum ensures that the cloud forms a flattened disk—theaccretion disk.
A few hundred thousand years after theBig Bang, theUniverse cooled to the point where atoms could form. As the Universe continued toexpand and cool, the atoms lost enough kinetic energy, anddark matter coalesced sufficiently, to formprotogalaxies. As further accretion occurred,galaxies formed.[7] Indirect evidence is widespread.[7] Galaxies grow throughmergers and smooth gas accretion. Accretion also occurs inside galaxies, forming stars.
Stars are thought to form insidegiant clouds of coldmolecular hydrogen—giant molecular clouds of roughly 300,000 M☉ and 65light-years (20 pc) in diameter.[8][9] Over millions of years, giant molecular clouds are prone tocollapse and fragmentation.[10] These fragments then form small, dense cores, which in turn collapse into stars.[9] The cores range in mass from a fraction to several times that of the Sun and are called protostellar (protosolar) nebulae.[8] They possess diameters of 2,000–20,000astronomical units (0.01–0.1 pc) and aparticle number density of roughly 10,000 to 100,000/cm3 (160,000 to 1,600,000/cu in). Compare it with the particle number density of the air at the sea level—2.8×1019/cm3 (4.6×1020/cu in).[9][11]
The initial collapse of a solar-mass protostellar nebula takes around 100,000 years.[8][9] Every nebula begins with a certain amount ofangular momentum. Gas in the central part of the nebula, with relatively low angular momentum, undergoes fast compression and forms a hothydrostatic (non-contracting) core containing a small fraction of the mass of the original nebula. This core forms the seed of what will become a star.[8] As the collapse continues, conservation of angular momentum dictates that the rotation of the infalling envelope accelerates, which eventually forms a disk.
As the infall of material from the disk continues, the envelope eventually becomes thin and transparent and theyoung stellar object (YSO) becomes observable, initially infar-infrared light and later in the visible.[11] Around this time the protostar begins tofusedeuterium. If the protostar is sufficiently massive (above80 MJ), hydrogen fusion follows. Otherwise, if its mass is too low, the object becomes abrown dwarf.[12] This birth of a new star occurs approximately 100,000 years after the collapse begins.[8] Objects at this stage are known as Class I protostars, which are also called youngT Tauri stars, evolved protostars, or young stellar objects. By this time, the forming star has already accreted much of its mass; the total mass of the disk and remaining envelope does not exceed 10–20% of the mass of the central YSO.[11]
At the next stage, the envelope completely disappears, having been gathered up by the disk, and the protostar becomes a classical T Tauri star.[13] The latter have accretion disks and continue to accrete hot gas, which manifests itself by strong emission lines in their spectrum. The former do not possess accretion disks. Classical T Tauri stars evolve into weakly lined T Tauri stars.[14] This happens after about 1 million years.[8] The mass of the disk around a classical T Tauri star is about 1–3% of the stellar mass, and it is accreted at a rate of 10−7 to 10−9 M☉ per year.[15] A pair of bipolar jets is usually present as well. The accretion explains all peculiar properties of classical T Tauri stars: strongflux in theemission lines (up to 100% of the intrinsicluminosity of the star),magnetic activity,photometricvariability and jets.[16] The emission lines actually form as the accreted gas hits the "surface" of the star, which happens around itsmagnetic poles.[16] The jets are byproducts of accretion: they carry away excessive angular momentum. The classical T Tauri stage lasts about 10 million years[8] (there are only a few examples of so-calledPeter Pan disks, where the accretion continues to persist for much longer periods, sometimes lasting for more than 40 million years[17]). The disk eventually disappears due to accretion onto the central star, planet formation, ejection by jets, andphotoevaporation byultraviolet radiation from the central star and nearby stars.[18] As a result, the young star becomes aweakly lined T Tauri star, which, over hundreds of millions of years, evolves into an ordinary Sun-like star, dependent on its initial mass.
Self-accretion ofcosmic dust accelerates the growth of the particles into boulder-sizedplanetesimals. The more massive planetesimals accrete some smaller ones, while others shatter in collisions. Accretion disks are common around smaller stars, stellar remnants in aclose binary, orblack holes surrounded by material (such as those at the centers ofgalaxies). Some dynamics in the disk, such asdynamical friction, are necessary to allow orbiting gas to loseangular momentum and fall onto the central massive object. Occasionally, this can result instellar surface fusion (seeBondi accretion).
In the formation of terrestrial planets orplanetary cores, several stages can be considered. First, when gas and dust grains collide, they agglomerate by microphysical processes likevan der Waals forces andelectromagnetic forces, forming micrometer-sized particles. During this stage, accumulation mechanisms are largely non-gravitational in nature.[19] However, planetesimal formation in the centimeter-to-meter range is not well understood, and no convincing explanation is offered as to why such grains would accumulate rather than simply rebound.[19]: 341 In particular, it is still not clear how these objects grow to become 0.1–1 km (0.06–0.6 mi) sized planetesimals;[5][20] this problem is known as the "meter size barrier":[21][22] As dust particles grow by coagulation, they acquire increasingly large relative velocities with respect to other particles in their vicinity, as well as a systematic inward drift velocity, that leads to destructive collisions, and thereby limit the growth of the aggregates to some maximum size.[23] Ward (1996) suggests that when slow moving grains collide, the very low, yet non-zero, gravity of colliding grains impedes their escape.[19]: 341 It is also thought that grain fragmentation plays an important role replenishing small grains and keeping the disk thick, but also in maintaining a relatively high abundance of solids of all sizes.[23]
A number of mechanisms have been proposed for crossing the 'meter-sized' barrier. Local concentrations of pebbles may form, which then gravitationally collapse into planetesimals the size of large asteroids. These concentrations can occur passively due to the structure of the gas disk, for example, between eddies, at pressure bumps, at the edge of a gap created by a giant planet, or at the boundaries of turbulent regions of the disk.[24] Or, the particles may take an active role in their concentration via a feedback mechanism referred to as astreaming instability. In a streaming instability the interaction between the solids and the gas in the protoplanetary disk results in the growth of local concentrations, as new particles accumulate in the wake of small concentrations, causing them to grow into massive filaments.[24] Alternatively, if the grains that form due to the agglomeration of dust are highly porous their growth may continue until they become large enough to collapse due to their own gravity. The low density of these objects allows them to remain strongly coupled with the gas, thereby avoiding high velocity collisions which could result in their erosion or fragmentation.[25]
Grains eventually stick together to form mountain-size (or larger) bodies called planetesimals. Collisions andgravitational interactions between planetesimals combine to produce Moon-size planetary embryos (protoplanets) over roughly 0.1–1 million years. Finally, the planetary embryos collide to form planets over 10–100 million years.[20] The planetesimals are massive enough that mutual gravitational interactions are significant enough to be taken into account when computing their evolution.[5] Growth is aided by orbital decay of smaller bodies due to gas drag, which prevents them from being stranded between orbits of the embryos.[26][27] Further collisions and accumulation lead to terrestrial planets or the core of giant planets.
If the planetesimals formed via the gravitational collapse of local concentrations of pebbles, their growth into planetary embryos and the cores of giant planets is dominated by the further accretions of pebbles.Pebble accretion is aided by the gas drag felt by objects as they accelerate toward a massive body. Gas drag slows the pebbles below the escape velocity of the massive body causing them to spiral toward and to be accreted by it. Pebble accretion may accelerate the formation of planets by a factor of 1000 compared to the accretion of planetesimals, allowing giant planets to form before the dissipation of the gas disk.[28][29] However, core growth via pebble accretion appears incompatible with the final masses and compositions ofUranus andNeptune.[30] Direct calculations indicate that, in a typicalprotoplanetary disk, the formation time of a giant planet via pebble accretion is comparable to the formation times resulting from planetesimal accretion.[31]
The formation ofterrestrial planets differs from that of giant gas planets, also calledJovian planets. The particles that make up the terrestrial planets are made from metal and rock that condensed in the innerSolar System. However, Jovian planets began as large, icy planetesimals, which then captured hydrogen and helium gas from thesolar nebula.[32] Differentiation between these two classes of planetesimals arise due to thefrost line of the solar nebula.[33]
Meteorites contain a record of accretion and impacts during all stages ofasteroid origin and evolution; however, the mechanism of asteroid accretion and growth is not well understood.[34] Evidence suggests the main growth of asteroids can result from gas-assisted accretion ofchondrules, which are millimeter-sized spherules that form as molten (or partially molten) droplets in space before being accreted to their parent asteroids.[34] In the inner Solar System, chondrules appear to have been crucial for initiating accretion.[35] The tiny mass of asteroids may be partly due to inefficient chondrule formation beyond 2AU, or less-efficient delivery of chondrules from near the protostar.[35] Also, impacts controlled the formation and destruction of asteroids, and are thought to be a major factor in their geological evolution.[35]
Chondrules, metal grains, and other components likely formed in thesolar nebula. These accreted together to form parent asteroids. Some of these bodies subsequently melted, formingmetallic cores andolivine-richmantles; others were aqueously altered.[35] After the asteroids had cooled, they were eroded by impacts for 4.5 billion years, or disrupted.[36]
For accretion to occur, impact velocities must be less than about twice the escape velocity, which is about 140 m/s (460 ft/s) for a 100 km (60 mi) radius asteroid.[35] Simple models for accretion in theasteroid belt generally assume micrometer-sized dust grains sticking together and settling to the midplane of the nebula to form a dense layer of dust, which, because of gravitational forces, was converted into a disk of kilometer-sized planetesimals. But, several arguments[which?] suggest that asteroids may not have accreted this way.[35]
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Comets, or their precursors, formed in the outer Solar System, possibly millions of years before planet formation.[37] How and when comets formed is debated, with distinct implications for Solar System formation, dynamics, and geology. Three-dimensional computer simulations indicate the major structural features observed oncometary nuclei can be explained by pairwise low velocity accretion of weak cometesimals.[38][39] The currently favored formation mechanism is that of thenebular hypothesis, which states that comets are probably a remnant of the original planetesimal "building blocks" from which the planets grew.[40][41][42]
Astronomers think that comets originate in both theOort cloud and thescattered disk.[43] The scattered disk was created whenNeptune migrated outward into the proto-Kuiper belt, which at the time was much closer to the Sun, and left in its wake a population of dynamically stable objects that could never be affected by its orbit (theKuiper belt proper), and a population whose perihelia are close enough that Neptune can still disturb them as it travels around the Sun (the scattered disk). Because the scattered disk is dynamically active and the Kuiper belt relatively dynamically stable, the scattered disk is now seen as the most likely point of origin for periodic comets.[43] The classic Oort cloud theory states that the Oort cloud, a sphere measuring about 50,000 AU (0.24 pc) in radius, formed at the same time as the solar nebula and occasionally releases comets into the inner Solar System as a giant planet or star passes nearby and causes gravitational disruptions.[44] Examples of such comet clouds may already have been seen in theHelix Nebula.[45]
TheRosetta mission to comet67P/Churyumov–Gerasimenko determined in 2015 that when Sun's heat penetrates the surface, it triggers evaporation (sublimation) of buried ice. While some of the resulting water vapour may escape from the nucleus, 80% of it recondenses in layers beneath the surface.[46] This observation implies that the thin ice-rich layers exposed close to the surface may be a consequence of cometary activity and evolution, and that global layering does not necessarily occur early in the comet's formation history.[46][47] While most scientists thought that all the evidence indicated that the structure of nuclei of comets is processedrubble piles of smaller ice planetesimals of a previous generation,[48] theRosetta mission confirmed the idea that comets are "rubble piles" of disparate material.[49][50] Comets appear to have formed as ~100-km bodies, then overwhelmingly ground/recontacted into their present states.[51]
