Nebular hypothesis

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In cosmogony, the nebular hypothesis is the currently accepted argument about how Earth's Solar System formed. It was first proposed in 1734 by Emanuel Swedenborg^[1]. In 1755 Immanuel Kant, who was familiar with Swedenborg's work, developed the theory further. He argued that nebulae slowly rotate, gradually collapsing and flattening due to gravity and eventually forming stars and planets, in a process of planetary formation. A similar model was proposed in 1796 by Pierre-Simon Laplace. These can be considered early theories of cosmology.

While originally applied only to our own Solar System, this method of planetary system formation was subsequently believed by theorists to be at work throughout the universe; over 250 extrasolar planets have since been discovered in our galaxy.

[edit] Overview of the solar nebula hypothesis

[edit] The original nebula

The hypothesis maintains that a planetary system begins as a large (typically ~10,000 AU in diameter), roughly spherical cloud of very cold interstellar gas, part of a larger molecular cloud. Such a nebula is just dense enough to begin contracting under the force of its own gravity, and its collapse may have been initiated by a pressure wave from a nearby event (such as a shock wave from a supernova) compressing the molecular cloud. The composition of such a nebula will reflect the composition of the resulting star; for our own Solar System the Solar Nebula is believed to have been comprised of about 98% (by mass) hydrogen and helium present since the Big Bang, and 2% heavier elements created by earlier generations of stars which died and ejected them back into interstellar space (see nucleosynthesis). The fraction of heavier elements is known as the cloud's metallicity; statistically, stars with larger metallicities (i.e. that formed from a cloud with more heavy elements) are more likely to possess planets. Once begun, the gravitational contraction of the solar nebula accelerates slowly but inevitably.

As it collapses, three physical processes shape the nebula: it heats up, its spin increases, and it flattens. The nebula heats up because atoms move more quickly as they fall deeper into the gravitational well and become denser, colliding more frequently: gravitational potential energy is converted to kinetic energy of the atoms, or thermal energy. Second, while initially imperceptible, the solar nebula had some small amount of net rotation (angular momentum), and because angular momentum is conserved, the nebula must rotate more quickly as it shrinks in size. Finally, the nebula must also flatten into a disk, called a protoplanetary disk, as collisions and mergers of blobs of gas average out their motions in favor of the direction of the net angular momentum.

Recent analysis of the composition of eight meteorites known to have formed at various times within the first three million years of the Solar system's formation has shown that around 1 million to 2 million years after the formation of the Sun the Solar system received an influx of Iron-60, thought to have been produced by the nearby supernova of a massive short-lived star produced in the same star-forming region as the Sun.^[2]

[edit] The protostar

Main article: protostar

An increasingly dense protostar accumulates at the solar nebula's center of gravity. During the process of planet formation in the disk, the protostar gradually compacts further, until after about 10 million to 50 million years, it finally reaches the conditions of temperature and pressure needed to initiate hydrogen nuclear fusion, and a star is born. A young star of this kind (a T Tauri star) produces a stellar wind, much stronger than that of a fully formed star, which eventually blows the remaining gases out of the disk, and largely ending the accretion process (particularly for any gas giants). Like most processes in a star's life, the time spent in the protostar phase depends on mass: massive stars collapse more quickly.

The gas in the protoplanetary disk, meanwhile, gradually cools from the gravitational heating of its collapse, and as it cools, dust (metals and silicates) and ice (hydrogen compounds such as water, methane, and ammonia) grains condense out of the gas (solidify). These grains gently bump into neighboring grains (collide) and stick together electrostatically, beginning the accretion process. Gas atoms and molecules are present in great abundance, but cannot be accreted, because they are moving too quickly to be held electrostatically. Hydrogen and helium, 98% of the mass of the disk, remain gaseous throughout the solar nebula, never condensing.

[edit] Planetesimals

Main article: planetesimal

The solid component of the disk is initially in the form of microscopic dust grains that seeded the precursor cloud; such grains in the interstellar medium are typically less than a micron in diameter, but through collisions in the protoplanetary disk they stick together and grow in size to become planetesimals (literally meaning an infinitely small planet). This dust is initially spread throughout the disk, but is expected to rain out into the disk midplane: just as the initial molecular cloud collapses under gravity into a disk, so the grains sink to the midplane but cannot move radially towards the protostar without losing angular momentum. Dust grains of different sizes fall down at different speeds, gathering more dust along the way.^[3] Larger grains may grow faster by clumping together randomly to produce fractal structures;^[4] such arrangements have more surface area for other grains to bump against and stick to. A population of large, fluffy grains may also inhibit the effects of gas drag,^[5] which could otherwise cause the solids to migrate towards the new-formed star before planets can form. This is called the core-accretion theory of planetary formation. ^[6] Fast collisions may shatter forming planetesimals, meaning the transition from dust to planetesimal is reversible. Turbulence in the disk may play a role in these collisions: if the turbulence is too violent, rainout into the midplane may be hindered, and destructive collisions between particles may be more common. Once planetesimals become sufficiently massive, their gravity helps bring more grains into contact,^[7] yet strong turbulence may also prevent this gravitational clumping, leading to growth through binary collision only. Nontheless, if gas giants are to form then planetesimals of about 1 km across must form within around 10,000 years.^[8]

Because planetesimals are so numerous, and spread throughout the protoplanetary disk, many survive the formation of a planetary system. Asteroids are understood to be left-over planetesimals, gradually grinding each other down into smaller and smaller bits, while comets are typically planetesimals from the farther reaches of a planetary system. Meteorites are samples of planetesimals that reach a planetary surface, and provide a great deal of information about the formation of our solar system. Primitive-type meteorites are chunks of shattered low-mass planetesimals, where no thermal differentiation took place, while processed-type meteorites are chunks from shattered massive planetesimals. Only the largest planetesimals survive these high-energy collisions with lower-mass planetesimals, and can continue to grow.

[edit] Oligarchic growth

As the planetesimals grow, they decrease in number and collisions become less frequent. Due to the stochastic nature of growth, not all planetesimals grow at the same rate and some will become more massive than others. As the planetesimals orbit the new star, dynamical friction keeps the energies (momentum) of the planetesimal population evenly distributed, so that the largest bodies possess the slowest velocities, typically orbiting in near-circular fashion, while the smaller planetesimals move much faster, on more eccentric orbits. Significantly, slower-moving bodies possess a larger collisional cross-section, the radius over which gravity can enhance a planetesimal's ability to capture another planetesimal. Consequently, the slower, more massive bodies are more effectively able to accrete the surrounding planetesimals, while the faster-moving, smaller bodies hardly grow at all. This quickly leads to a runaway process, where the largest bodies in each region of the disk come to dominate, growing much larger than the surrounding "planetesimal sea".^[9][10] These massive bodies now completely dominate the solid material in the disk and are called oligarchs, meaning the few that rule; the process is known as oligarchic growth.^[11] These few planetesimals rapidly increase in size, increasing from a few tens of kilometres across before oligarchic growth begins, to several hundred and eventually several thousand kilometres in diameter.

The process of oligarchic growth is self-limiting: each oligarch has a limited feeding zone (determined by its collisional cross-section), and can grow no further once all the planetesimals within it have been accreted. It is doubtful whether these zones contain enough solids for the oligarchs to grow to terrestrial masses, suggesting the planetesimals' growth should stall at a few hundred kilometres in size.^[12] However, it is possible that turbulence once again plays a part, as it can deposit or take away angular momentum from the planetesimals, providing a random component of radial motion. This may provide a steady influx of new material to the feeding zones, allowing the oligarchs to continue with their growth.^[13]

One way or another the oligarchs continue to grow, until (interior to the frost line) they have reached masses of typically 0.05-1 Earth mass within a million years or so,^[14] large enough to be considered protoplanets; bodies in the outer disk will grow larger still, due to the higher density of solids available for growth. The terrestrial planetary region now likely consists of a few dozen widely spaced oligarchs,^[11] dynamically isolated and unlikely to collide for hundreds of thousands or even millions of years.

[edit] Non-uniform temperatures

The temperature in a protoplanetary disk is not uniform, and this is the key to understanding the differentiation between terrestrial and jovian planet formation. Inside the frost line, the temperature is too high (above 150 K) for hydrogen compounds to condense: they remain gaseous. The only grains available for accretion, then, are the heavier metal and silicate dust grains. Thus the planetesimals in this region are composed entirely of rock and metal, such as the asteroids, and make up the terrestrial planets.

Outside the frost line, hydrogen compounds such as water, methane and ammonia are able to solidify into 'ice' grains, and accrete. Rock and metal grains are also available, but are vastly outnumbered (and outweighed) by the hydrogen compounds, which are much more abundant everywhere. Thus the planetesimals in this region are icy bodies with small amounts of rock and metal mixed in. The Kuiper Belt and Oort Cloud objects, comets, Neptune's huge moon Triton, and probably Pluto and its moon Charon, are all examples of these 'dirty snowball' planetesimals. Due to the greater amount of solid materials available, as well as less frequent collisions and lower velocities (being in much larger orbits), the largest of these planetesimals grow so massive (about 10 times the mass of the Earth) their gravity begins to collect and retain helium and even hydrogen gases. Once that starts, they grow rapidly, as hydrogen and helium are 98% of the disk, and collecting these gases increases their mass and consequently the size of their gravitational net.

[edit] Jovian planetesimals

Soon the jovian planetesimals are nothing like the icy bodies they came from, but are more or less dominated by the hydrogen and helium gas they have captured, huge gaseous clouds with dense cores. These jovian gas balls then, in close analogy to the solar system itself, gradually collapse gravitationally, heating up, rotating more quickly, and flattening. Some moons of the jovian planets may be formed in an analogous process to the planets themselves, coalescing from condensed grains in the disks which formed as the gas giant protoplanet collapsed. This can explain why, in our own Solar System, the jovian planets all have many moons and rings in the same plane, and why jovian planets rotate quickly. The growth of the jovians ends when the young star's strong stellar wind blows the remaining gas and dust out of the disk and into interstellar space.

In the simplest possible terms, the innermost giant protoplanetary core forms where the disk density is highest and dynamical times (the typical timescale for collisions) are shortest; hence this body reaches the critical mass for gas capture earliest, and in the densest regions of the disk, and so has longest to accrete the surrounding gas. In our own Solar System Jupiter was the largest protoplanetary core beyond the frost line, and so fulfilled this role, becoming the largest planet in that system. In reality, the process may be more complicated, with planetary migration and turbulence muddying this picture; compared to the extra solar planetary systems observed to date, the distribution of the planets in our own system may even be considered somewhat unusual.

[edit] Giant impacts

Finally, after the stellar wind has cleared the gas out of the disk, a large population of protoplanets and planetesimals may be left over. Over a period of 10 million to 100 million years, these protoplanets — typically with a mass between that of the Moon and several Earths — will perturb each other until orbit-crossing occurs, leading to collisions. The bodies resulting from these collisions will be the final planets of that system. Such a collision, between the proto-Earth and a Mars-sized protoplanet, is believed to have formed the Earth's moon. The process is highly random; a forming terrestrial system near-identical to that which produced the inner planets may easily end with fewer or more planets than we observe in our Solar System.

The smaller planetesimals, being vastly more numerous, remain within the planetary system for much longer. These bodies may be swept up by the planets that have formed in a process known as "clearing the neighbourhood", either by slinging them in the distant outer reaches of the system (the Oort Cloud in our Solar System), or continually nudging their orbits into collisions or stable orbits with other planets. This period of bombardment lasts several hundred million years, and may leave evidence of cratering which is still visible on geologically dead bodies. In some respects, as long as there are small rocky or icy bodies available to the system, it may be argued that this stage of formation never really 'finishes', as the threat of asteroid collisions with Earth or the recent impact of comet Shoemaker-Levy 9 upon Jupiter ably demonstrates.

In our own Solar System it is believed this episode of formation was exceptionally strong due to a 2:1 resonance crossing of Jupiter and Saturn, catastrophically disturbing a large outer planetesimal disk, and the process is known as the Late Heavy Bombardment.

[edit] Solar system features explained by theory

The nebular hypothesis effectively explains all the major features of our solar system:

1. Regular motions of the planets and moons (all revolve in the nearly same plane, in nearly circular orbits, in same direction the Sun rotates, and nearly all rotate in the nearly same direction too)
2. All major differences between terrestrial and jovian planets (mass, distance from Sun, composition, moon and ring systems)
3. Small bodies (asteroids and comets, both short- and long-period)
4. Exceptions to the trends (terrestrial moons, axial tilts, non-coplanar jovian moons, Triton)

[edit] Challenges to the hypothesis

The current challenges for the nebular hypothesis include explaining:

1. missing mass in Kuiper Belt
2. capture process for Triton
3. sideways rotation of Uranus
4. discovered "hot Jupiter" exoplanets
5. discovered exoplanets in binary and trinary stellar systems

[edit] The meaning of accretion

Use of the term accretion disk for the protoplanetary disk leads to confusion over the planetary accretion process.

The protoplanetary disk is sometimes referred to as an accretion disk, because while the young T Tauri-like protosun is still contracting, gaseous material may still be falling onto, accreting on, its surface from the disk's inner edge.

However, that meaning should not to be confused with the process of accretion forming the planets. In this context, accretion refers to the process of cooled, solidified grains of dust and ice orbiting the protosun in the protoplanetary disk, colliding and sticking together and gradually growing, up to and including the high energy collisions between sizable planetesimals.

In addition, the jovians probably had accretion disks of their own, in the first meaning of the word. The clouds of captured hydrogen and helium gas contract, spin up, flatten, and deposit gas onto the surface of each jovian protoplanet, while solid grains within that disk accrete into planetesimals and eventually forming the jovian moons.

[edit] See also

• Formation and evolution of the Solar System
• History of Earth
• Asteroid Belt, Kuiper Belt, and Oort Cloud
• Bok globule, Herbig-Haro object
• T Tauri star

[edit] References

1. ^ Swedenborg, Emanuel. 1734, (Principia) Latin: Opera Philosophica et Mineralia (English: Philosophical and Mineralogical Works), (Principia, Volume I)
2. ^ Bizzarro, Martin; David Ulfbeck, Anne Trinquier, Kristine Thrane, James N. Connelly, Bradley S. Meyer (25 May 2007). "Evidence for a Late Supernova Injection of ⁶⁰Fe into the Protoplanetary Disk". Science 316 (5828): 1178 - 1181. doi:10.1126/science.1141040.
3. ^ Weidenschilling S.J., (1980). "Dust to planetesimals - Settling and coagulation in the solar nebula". Icarus 44: 172-189.
4. ^ Meakin P.; Donn B., (1988). "Aerodynamic properties of fractal grains - Implications for the primordial solar nebula". Astrophysical Journal, Part 2 - Letters 329: L39-L41.
5. ^ Takeuchi T.; Clarke C. J.; Lin D. N. C., (2005). "The Differential Lifetimes of Protostellar Gas and Dust Disks". The Astrophysical Journal 627: 286-292.
6. ^ Laughlin, Gregory P. (2006), "Extra solar Planetary Systems", American Scientist 94 (5): pp. 420-429
7. ^ Goldreich P.; Ward W.R., (1973). "The Formation of Planetesimals". Astrophysical Journal 183: 1051-1062.
8. ^ Lissauer J. J., (1993). "Planet formation". Annual review of astronomy and astrophysics 31: 129-174.
9. ^ Wetherill G. W.; Stewart G. R., (1989). "Accumulation of a swarm of small planetesimals". Icarus 77: 330.
10. ^ Ohtsuki K.; Ida S., (1990). "Runaway planetary growth with collision rate in the solar gravitational field". Icarus 85: 499-511.
11. ^ ^{a b} Kokubo E.; Ida S., (2000). "Formation of Protoplanets from Planetesimals in the Solar Nebula". Icarus 143: 15-27.
12. ^ Lissauer J. J., (1987). "Timescales for planetary accretion and the structure of the protoplanetary disk". Icarus 69: 249-265.
13. ^ Fogg M. J.; Nelson R. P., (2005). "Oligarchic and giant impact growth of terrestrial planets in the presence of gas giant planet migration". Astronomy & Astrophysics 441: 791-806.
14. ^ Weidenschilling S. J.; Spaute D.; Davis D. R.; Marzari F.; Ohtsuki K., (1997). "Accretional Evolution of a Planetesimal Swarm". Icarus 128: 429-455.

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