Talk:Uranus/new subpage
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Note: Nebula : Noun.
- Nebular: Adjective
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[edit] The formation of the stars and disks
Stars are thought to form inside giant clouds of cold molecular hydrogen roughly 300,000 times the mass of the Sun and 20 parsecs in diameter.[1][2] Over millions of years, giant molecular clouds are prone to collapse and fragmentation.[3] These fragments then form small, dense cores which in turn collapse into stars.[1] These cores range in mass from a fraction to several times that of the Sun. They possess diameters of 0.01-01 pc (2000-20,000 AU) and a core particle number density of roughly 10,000 to 100,000 cm-3.[4][1][5] These cores are called protostellar nebulae (protosolar in the case of Sun).[2]
The initial collapse of a solar-mass protostellar nebula takes of the order of 100,000 years.[1][2] Every nebula begins its collapse with a certain amount of angular momentum. Gas in the central part of the nebula, whose angular momentum is relatively low, undergoes fast compression and forms a hot hydrostatic (not contracting) core with a fraction of the mass of the original nebula.[6] This core forms the seed of what will become a star.[6][2] As the collapse continues, conservation of angular momentum means the rotation of the infalling gas accelerates.[7][8] This rotation largely prevents the gas from directly accreting onto the central core and instead forces it to spread outwards near its equatorial plane, forming a disk that in turn accretes onto the central core.[2][7][8] The core gradually grows in mass until it becomes a hot protostar.[6] At this stage, unlike more evolved protostars, the younger protostar and its disk are heavily obscured by the infalling envelope and are not directly observable.[9] In fact the remaining envelope's opacity is so high that even millimeter-wave radiation can not escape from inside it.[9][2] Such objects are observed as very bright condensations, which emit mainly millimeter-wave and submilimiter-wave radiation.[5] They are classified as (spectral) Class 0 protostars.[9] However the collapse is often accompanied by bipolar outflows (jets), which emanate along the rotational axis of the inferred disk. Such outflows are often observed in star-forming regions (see Herbig-Haro objects).[10] At this stage the protostar has a high luminosity in the millimeter-wave and submilimiter-wave spectral regions— a protostar of solar mass may radiate at up to one hundred solar luminosities— and does not fuse hydrogen.[6] Their main source of energy is gravitational collapse.[6]
As the envelope's material falls into the disk, it eventually becomes thin and transparent and the young stellar object (YSO) becomes observable; initially in far-infrared light and later in the visible.[5] Around this time the protostar also begins to fuse hydrogen.[6] This event, which can be called the birth of a new star, happens at approximately 100,000 years after the collapse began.[2] The external appearance of the YSO at this stage corresponds to the spectral class I protostars,[9] also called young T Tauri stars or evolved protostars.[9] By this stage 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.[5]
At the next stage of evolution, the envelope disappears, having been accreted by the disk, and the star (or protostar) becomes a classical T Tauri star. This happens after about 1 million years.[2] A T Tauri star is a young solar mass star with hightened levels of stellar activity. T Tauri stars are divided into two classes: weakly lined and classical. Weakly lined T Tauri stars do not possess accretion disks; in contrast, classical T Tauri stars have accretion disks and continue to accrete hot gas, which manifests as strong emmision lines in the spectrum. Classical T Tauri stars evolve into weakly lined T Tauri stars. The mass of the remaining disk is still an appreciable fraction of the stellar mass, and it is accreted onto the star at the rate of between a 10 milionth to one billionth a solar mass per year.[11] A pair of bipolar jets is usually present as well.[12] The accretion explains all peculiar properties of the classical T Tauri stars: strong flux in the emission lines (up to 100% of the intrinsic luminosity of the star), magnetic activity, photometric variability and jets.[13] The emission lines actually form as the accreted gas hits the "surface" of the star, which happens around its magnetic poles.[13] The jets are byproducts of accretion: they carry away excessive angular momentum. The classical T Tauri stage lasts about 10 million years.[2] The disk eventually disappears due to accretion onto central star, planet formation, ejection by jets and photoevaporation by UV-radiation from the cental star and nearby stars.[14] As a result the young star becomes a weakly lined T Tauri star, which slowly, over timeframe of hundreds of millions of years, evolves into an ordinary sun-like star.[6] I think you should elaborate on this. Surely the process by which a star becomes ordinary should be discussed?
[edit] Disk evolution
Under some circumstances, the disk, which can now be called protoplanetary, gives birth to a planetary system.[2] As was shown above, protoplanetary disks are ubiquitous around all Sun-like stars.[15][16] They exist from the beginning of a star's formation, but at the earliest stages are unobservable due to the opacity of the surrounding envelope.[9] At the earliest stage (Class 0 protostar) the disk is thought to be massive and hot. It is an accretion disk, which feeds the central protostar.[7][8] The temperature can easily exceed 400 K inside 5 AU and 1000 K inside 1 AU.[17] The heating of the disk is primarily caused by the vicious dissipation of turbulence in it and by the infall of the gas from the nebula.[7][8] The high temperature in the inner disk causes most of the volatile material (water, organics, and even some rock) to evaporate, leaving only the most refractory elements like iron. Volatile dust can survive only in the outer part of the disk.[17]
The main problem in the physics of accretionary disks is the generation of turbulence and the mechanism responsible for the high effective viscosity.[2] Such a turbulent viscosity is thought to be responsible for the transport of the mass to the central protostar and momentum to the periphery of the disk. Such a transport is vital for accretion, because gas can be accreted by the central protostar only if it losses most of its angular momentum.[7] Since the momentum is conserved, a part of the gas must drift outward carrying the excesive momentum.[7][18] The result of this transport is the growth of both the protostar and of the disk radius, which can reach 1000 AU if the initial angular momentum of the nebula is large enough.[8] Such large disks are routinely observed in many star-forming regions,[19] such as the Orion nebula.
The lifespan of the disk is about 100 million years.[16] By the time the star reaches the classical T-Tauri stage, the disk becomes thinner[11] and cools. Less volatile materials start to condense in the inner disk forming dust grains. These grains have size of 0.1-1 μm and contain crystalline silicates.[20] The transport of the material from the outer disk can mix these newly formed dust grains with primordial ones, which contain organic matter and other volatiles. Such a mixing can explain some peculiarities in the composition of solar system bodies i.e. presence of interstellar grains in the primitive meteorites. In the dense disk environment dust particles tend to stick to each other, leading to the formation of larger particles up to several centimeters in size. The signatures of the dust processing and coagulation are observed in the infrared spectra of the young disks.[20] Further evolution of the dust particles can lead to the formation of planetesimals with a size of 1 km or larger, which are building blocks of planets.[2] Planetesimal formation is another unsolved problem of disk physics, because simple sticking becomes ineffective as dust particles grow larger.[21] The favorite hypothesis is formation by the gravitational instability. Particles several ctntimeters in size or larger slowly settle near the middle plane of the disk. This middle layer can become very thin (less than 100 km) and dense. Such a layer is gravitationally unstable and will fragment into numerous clumps, which will collapse into planetasimals.[21][2]
Planetary formation can also be triggered by gravitational instability within the disk itself, which leads to its fragmentation into clumps. Some of these clumps, if they dense enough, can collapse, forming gas giant planets.[18] Gravitational instability can cause rapid formation of gas giant planets and even brown dwarfs at the timescale of 1000 years.[22] However it is only possible in massive disks (more massive than 0.3 solar masses). In comparizon typical disk masses are 0.01–0.1 solar masses. Because such massive disks are rare, this mechanism of the planet formation is thought to be infrequent.[23][2]
The ultimate dissipation of protoplanetary disks is triggered by a number of different mechanisms. The inner part of the disk is either accreted by the star or ejected by the bipolar jets,[11][12] whereas the outer part can evaporate under the star's UV radiation during its T Tauri stage.[24] or nearby stars.[14] The gas in the central part can either be accreted or ejected by the growing planets, while the small dust particles are ejected by the radiation pressure of the central star. What is finally left is either a planetary system, a remnant disk of dust without planets, or nothing, if planetesimals failed to form.[2]
[edit] Rocky planet formation
Rocky planets form in the inner part of the protoplanetary disk, where temperature is high enough to prevent condensation of water and other ices.[25] This results in coagulation of purely rocky grains and later in the formation of rocky planetesimals.[25] Such conditions are thought to exist in the inner 3–4 AU part of the disk of a sun-like star.[2]
After planetesimals (about 1 km in diameter) have formed by one way or another, runaway accretion begins.[26] It is called runaway because the mass growth rate is proportional to R4~M4/3, where R and M are the radius and mass of the growing body, respectively.[27] It is obvious that the specific (divided by mass) growth accelerates as the mass increases. Such accretion leads to preferential growth of larger bodies at the expense of smaller ones.[26] The runaway accretion lasts between ten and a hundred thousand years and ends when the largest bodies exceeds approximately 1000 km in diameter.[26] Slowing of the accretion is caused by gravitational perturbations by large bodies on the remaining planetesimals.[27][26] In addition, influences of large bodies at this stage prevents further growth of any small bodies.[26]
The next stage is called oligarchic accretion.[26] It is characterized by the dominance of several hundred of the largest bodies – oligarchs, which continue to slowly accrete planetesimals.[26] No bodies other than the oligarchs can grow.[27] At this stage the rate of accretion is proportional to R2, i.e. to the cross-section of an oligarch. [27] The specific accretion rate is proportional to M-1/3, i.e. it declines with the mass of the body. This allows smaller oligarchs to catch up to larger ones. The oligarchs are kept at the distance of about 10·Hr (Hr =(M/3M)1/3 is Hill radius) from each other by the influence of the remaining planetesimals.[26] Their orbital eccentricities and inclinations remain small. The oligarchs continue to accrete until planetesimals are exhausted in the disk around them.[26] The final mass of an oligarch depends on the distance from the star and surface density of planetesimals and is called the isolations mass.[27] For the rocky planets it is up to 0.1 of the Earth mass, or one Mars mass.[2] The final result of the oligarchic stage is the formation of about 50 planetary embryos uniformly spaced at about 10·Hr.[28] They are thought to reside inside gaps in the disk and to be separated by rings of remaining planetesimals. This stage is thought to last a few 105 years.[26][2]
The last stage of rocky planet formation is the merger stage.[2] It begins when only a small number of planetesimals remains and embryos become massive enough to perturb each other, which causes their orbits to become chaotic.[28] During this stage embryos expel remaining planetesimals, and collide with each other. The result of this process, which lasts for ten to a hundred million years, is the formation a limited number of Earth sized bodies. Simulation show that the number of surviving planets is on average from 2 to 5.[29][30][28][2] In the Solar System they may be represented by Earth and Venus.[28] Formation of both planets required merging of approximately 10 embryos, while the equal number of them were thrown out of the Solar System.[30] Some the embryos, which originated in the asteroid belt, are thought to have brought water to Earth.[25] Mars and Mercury may be regarded as remaining embryos that survived that rivalry.[30] Rocky planets that have managed to coalesce settle eventually into more or less stable orbits, explaining why planetary systems are generally packed to the limit; or in other words why they always appear to be at the brink of instability.[28]
[edit] Giant planet formation
The formation of giant planets is an outstanding problem in the planetary sciences.[23] In principle two possibilities exist. The first one is a disk instability model, where giant planets form in the massive protoplanetary disks as a result of its gravitational fragmentation (see above).[22] The disk instability may also lead to the formation of brown dwarfs, which are usually classified as stars. The second possibility is a core accretion model, which is also known as a nucleated instability model.[23] The latter scenario is thought to be the most promising one, because it can explain the formation of the giant planets in relatively low mass disks (less than 0.1 solar masses). In this model, giant planet formation is divided into two stages: a) Accretion of a core of approximately 10 Earth masses and b) accretion of gas from the protoplanetary disk.[23][2]
Giant planet core formation is thought to proceed roughly along the lines of terrestrial planet formation.[26] It starts with planetesimals, which then undergo the runaway growth followed by the slower oligarchic stage.[27] Hypotheses do not predict a merger stage, due to the low probability of collisions between planetary embryos in the outer parts of planetary systems.[27] An additional difference is the omposition of the planetesimals, which in the case of giant planets form beyond the so called snow line and consist mainly of ice (ice to rock ratio is about 4 to 1).[31] This enhances the mass of planetesimals four times. However the minimum mass nebular, which is capable of terrestrial planet formation, can only form 1-2 Earth mass cores at the distance of Jupiter (5 AU) within ten million years.[27] The latter number represents an averages lifetime of gaseous disks around sun-like stars.[16] The proposed solutions include enhanced mass of the disk (a tenfold increase would suffice);[27] protoplanet migration, which allows the embryo to accrete more planetesimals;[31] and finally accretion enhancement due to gas drag in the gaseous envelopes of the embryos.[31][32] Some combination of the above-mentioned ideas may explain the formation of the cores of gas giant planets such as Jupiter and perhaps even Saturn.[23] The formation of planets like Uranus and Neptune is more problematic, since no theory has been capable of providing for the in situ formation of their cores at the distances of 20-30 AU from the central star.[2] To resolve this issue an idea has been brought forward that they initially accreted in the Jupiter-Saturn region and then were scattered and migrated to their present location.[33]
Once the cores are of sufficient mass, they begin to gather gas from the surrounding disk.[2] Initially it is a slow process, which can increase the core masses up to 30 earth masses in a few million years.[31][32] After that the accretion rates increase dramatically and the remaining 90% of the mass is accumulated in 104 years.[32] The accretion of the gas stops, when it is exhausted, i.e. when a gap opens in the protoplanetary disk.[34] In this model ice giants – Uranus and Neptune are failed cores that began gas accretion too late, when almost all gas had already disappeared. The post runaway gas accretion stage is characterized by migration of the newly formed giant planets and continued slow gas accretion.[34] Migration is caused by the interaction of the planet sitting in the gap with the remaining disk. It stops, when the protoplanetary disk disappears or when the end of the disk is attained. The latter case corresponds to the so called hot Jupiter, which are likely to have stopped their migration, when they reached inner hole in the protoplanetary disk.[34]
Giant planets can significantly influence terrestrial planet formation. The presence of giants tends to increase eccentricities and inclinations of planetesimals and embryos in the terrestrial planet region (inside 4 AU in the Solar System).[30][29] On the one hand, if giant planets form too early they can slow or prevent inner planet accretion. On the other hand, if giant planets form near the end of the oligarchic stage, as is thought to have happened in the Solar System, they will influence the merges of planetary embryos making them more violent.[30] As a result the number of terrestrial planets will decrease and they will be more massive.[35] In addition, the size of the system will shrink, i.e. terrestrial planets will form closer to the central star. In the Solar System the influence of giant planets (particularly of Jupiter) is thought to have been limited because they are relatively remote from the terrestrial planets (inside 2 AU).[35]
However the region of a planetary system, which is adjacent to the giant planets, will be influenced in a different way.[29] In such a region eccentricities of embryos may become so large that they may pass close to a giant planet. As a result they may (and probably will) be thrown out of the planetary system.[36][30][29] If all embryos are removed then no planets will form in this region.[29] Additional consequence is that a huge number of small planetesimals will remain, because giant planets along (without help from embryos) are incapable of clearing them all out. Although the total mass of remaining planetesimals will be small, because cumulative action of the embryos (before their ejection) and giant planets is still strong enough to remove 99% of the small bodies.[30] Such a region will eventually evolve in an asteroid belt, which is a full analog of the main asteroid belt in the Solar System (from 2 to 4 AU from the Sun).[30][29]
A few general notes:
You should clarify the difference between the two different kinds of "accretion" discussed in this article
The accretion can be defined generally as a process, by which small bodies or gas grow into larger bodies. This definition implies that there is only one accretion. Ruslik (talk) 07:16, 20 November 2007 (UTC)
Also, it might be a good idea to briefly explain what T-Tauri stars are, and why we think they are similar to the young Sun.
Added a note. Ruslik (talk) 07:24, 20 November 2007 (UTC)
[edit] What is known and unknown about planetary formation
Which is why the Known section really belongs up top. It would make a great introduction.
The position is currently not important, because it is only a draft version. The position in the real article will be determined later. Ruslik (talk) 07:13, 20 November 2007 (UTC)
[edit] Known
The star formation process naturally results in appearance of accretionary disks around young stellar objects.[9] At the age of about 1 million years 100% of stars may have such disks.[16] This conclusion is supported by observations of the gaseous and dusty disks around protostars and T Tauri stars as well as by theoretical considerations.[19] The observations of the disks show that the dust grains inside them grow in size on the short time scale producing 1 cm sized particles.[20]
The accretion process, by which 1 km planetesimals grow into 1000 km sized bodies, is well understood now.[26] This process develops inside any disk, where the number density of planetesimals is sufficiently high, and proceeds in a runaway manner. At the later stage the growth slows and continues as the oligarchic accretion. The end result is formation of planetary embryos with varying sizes, which depend on the distance from the star.[26]
Various simulations have demonstrated that the merger of thus embryos in the inner part of the planetary system leads to the formation of a few Earth sized bodies. So the origin of terrestrial planets is now considered to be an almost solved problem.[28]
[edit] Unknown
The physics of the accretionary disks is not well understood.[23] The main problem is the mechanism of angular momentum transport from the inner to the outer part of the disk, which is necessary for the efficient accretion. The process (or processes), which is responsible for the disappearance of the disks, are also not well known.[18][7]
The formation of planetesimals is a biggest unsolved problem in the theory of planet formation. The precise mechanism by which 1 cm particles coalesce into the 1 km planetesimals is not understood. This mechanism appears to be the key to the question why some stars have planets, while others have nothing around them (even dusty belts).[21]
The formation of giant planets is another unsolved problem. Current theories have a serious trouble explaining how their cores can form fast enough to accrete significant amounts of gas from the quickly disappearing protoplanetary disk.[26][31] The lifetimes of the disks (<107 years) appears to be shorter than timescales of the core formations.[16] Another problem of giant planet formation is their migration. Some calculations show that the interactions with the disk can cause the rapid inward migration, which if not stopped will result in them plunging into the star.[34]
[edit] References
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- ^ Compare it with the particle number density in air at sea level—2.8×1019 cm-3
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- ^ As a variant they may collide with the central star or a giant planet
