STARS

STARS

Star formation is the process by which dense parts of molecular clouds collapse into a ball of plasma to form a star. An interstellar cloud of gas will remain in hydrostatic equilibrium as long as the kinetic energy of the gas pressure is in balance with the potential energy of the internal gravitational force. If a cloud is massive enough that the gas pressure is insufficient to support it, the cloud will undergo gravitational collapse depending on the temperature and density of the cloud, but is typically thousands to tens of thousands of solar masses.

In triggered star formation, one of several events might occur to compress a molecular cloud and initiate its gravitational collapse. Molecular clouds may collide with each other, or a nearby supernova explosion can be a trigger, sending shocked matter into the cloud at very high speeds. Alternatively, galactic collisions can trigger massive starbursts of star formation as the gas clouds in each galaxy are compressed and agitated by tidal forces.

A super massive black hole at the core of a galaxy may serve to regulate the rate of star formation in a galactic nucleus.

Stellar evolution is the process by which a star undergoes a sequence of radical changes during its lifetime. Depending on the mass of the star, this lifetime ranges from only a few million years for the most massive to trillions of years for the least massive, which is considerably longer than the age of the universe.

All stars are born from collapsing clouds of gas and dust, often called nebulae or molecular clouds. Nuclear fusion powers a star for most of its life. Stars similar to our Sun gradually grow in size until they reach a red giant phase, after which the core collapses into a dense white dwarf and the outer layers are expelled as a planetary nebula. Larger stars can explode in a supernova as their cores collapse into an extremely dense neutron star or black hole. It is not clear how red dwarfs die because of their extremely long life spans, but they probably experience a gradual death in which their outer layers are expelled over time.

Stellar evolution begins with the gravitational collapse of a giant molecular cloud (GMC). Typical GMCs are roughly 100 light-years (9.5×10¹⁴ km) across and contain up to 6,000,000 solar masses (1.2×10³⁷ kg). As it collapses, a GMC breaks into smaller and smaller pieces. In each of these fragments, the collapsing gas releases gravitational potential energy as heat. As its temperature and pressure increase, a fragment condenses into a rotating sphere of super hot gas known as a protostar.

The further development heavily depends on the mass of the evolving protostar; in the following, the protostar mass is compared to the solar mass: 1.0 M_☉ (2.0×10³⁰ kg) means 1 solar mass. Protostars with masses less than roughly 0.08 M_☉ (1.6×10²⁹ kg) never reach temperatures high enough for nuclear fusion of hydrogen to begin. These are known as brown dwarfs. Brown dwarfs heavier than 13 Jupiter masses (2.5 × 10²⁸ kg) or 0.0125 solar mass fuse deuterium. Both types, deuterium-burning or not, shine dimly and die away slowly, cooling gradually over hundreds of millions of years.

For a more massive protostar, the core temperature will eventually reach 10 million kelvins, initiating the proton-proton chain reaction and allowing hydrogen to fuse, first to deuterium and then to helium. The onset of nuclear fusion leads relatively quickly to a hydrostatic equilibrium in which energy released by the core exerts a "radiation pressure" balancing the weight of the star's matter, preventing further gravitational collapse. The star thus evolves rapidly to a stable state, beginning the main sequence phase of its evolution.

Small, relatively cold, low mass red dwarfs burn hydrogen slowly and will remain on the main sequence for hundreds of billions of years, while massive, hot super giants will leave the main sequence after just a few million years. A mid-sized star like the Sun will remain on the main sequence for about 10 billion years. The Sun is thought to be in the middle of its lifespan.

The sun is like a middle aged man ! Half way through.

A mid-sized star like the Sun will remain on the main sequence for about 10 billion years. Eventually, the core exhausts its supply of hydrogen. Without the outward pressure generated by the fusion of hydrogen to counteract the force of gravity, it contracts until either electron degeneracy becomes sufficient to oppose gravity or the core becomes hot enough (around 100 mega kelvins) for helium fusion to begin. Which of these happens first depends upon the star's mass.

Low-mass stars

What happens after a low-mass star ceases to produce energy through fusion is not directly known: the universe is thought to be around 13.7 billion years old, which is less time than it takes for the fusion to cease in such stars.

Some stars may fuse helium in core hot-spots, causing an unstable and uneven reaction as well as a heavy solar wind. In this case, the star will form no planetary nebula but simply evaporate, leaving little more than a brown dwarf.

A star of less than about 0.5 solar mass will never be able to fuse helium even after the core ceases hydrogen fusion. There simply is not a stellar envelope massive enough to exert enough pressure on the core. These are the red dwarfs, such as Proxima Centauri, some of which will live thousands of times longer than the Sun. Recent astrophysical models suggest that red dwarfs of 0.1 solar mass may stay on the main sequence for almost six trillion years, and take several hundred billion more to slowly collapse into a white dwarf.

If a star's core becomes stagnant (as is thought will be the case for the Sun), it will still be surrounded by layers of hydrogen which the star may subsequently draw upon. However, if the star is fully convective (as thought to be the case for the lowest-mass stars), it will not have such surrounding layers. If it does, it will develop into a red giant as described for mid-sized stars below, but never fuse helium as they do; otherwise, it will simply contract until electron degeneracy pressure halts its collapse, thus directly turning into a white dwarf.

Mid-sized stars

Stars of roughly 0.5–10 solar masses become red giants of two types. Red giant branch stars (RGB stars) whose shells are still fusing hydrogen into helium, while the core is inactive helium. They have reached hydrostatic equilibrium, where electron degeneracy pressure is sufficient to balance gravitational pressure.

A star of up to a few solar masses will develop a helium core supported by electron degeneracy pressure, surrounded by layers which still contain hydrogen. Its gravity compresses the hydrogen in the layer immediately above it, causing it to fuse faster than hydrogen would fuse in a main-sequence star of the same mass. This in turn causes the star to become more luminous (from 1,000–10,000 times brighter) and expand; the degree of expansion outstrips the increase in luminosity, causing the effective temperature to decrease.

The star forms from a collapsing cloud of gas (1),
and then undergoes a contraction period as a protostar (2), before joining the main sequence (3). Once the Hydrogen at the core is consumed it expands into a red giant (4), then sheds its envelope into a planetary nebula and degenerates into a white dwarf (5).

As the hydrogen around the core is consumed, the core absorbs the resulting helium, causing it to contract further, which in turn causes the remaining hydrogen to fuse even faster. This eventually leads to ignition of helium fusion in the core. In stars of more than approximately 0.5 solar masses, electron degeneracy pressure may delay helium fusion for millions or tens of millions of years; in more massive stars, the combined weight of the helium core and the overlying layers means that such pressure is not sufficient to delay the process significantly.

Massive stars

In massive stars, the core is already large enough at the onset of hydrogen burning shell that helium ignition will occur before electron degeneracy pressure has a chance to become prevalent. Thus, when these stars expand and cool, they do not brighten as much as lower mass stars; however, they were much brighter than lower mass stars to begin with, and are thus still brighter than the red giants formed from less massive stars. These stars are unlikely to survive as red super giants; instead they will destroy themselves as type II supernovas.

Extremely massive stars (more than approximately 40 solar masses), which are very luminous and thus have very rapid stellar winds, lose mass so rapidly due to radiation pressure that they tend to strip off their own envelopes before they can expand to become red super giants, and thus retain extremely high surface temperatures (and blue-white color) from their main sequence time onwards. Stars cannot be more than about 120 solar masses because the outer layers would be expelled by the extreme radiation.

The core grows hotter and denser as it gains material from fusion of hydrogen at the base of the envelope. In all massive stars, electron degeneracy pressure is insufficient to halt collapse by itself, so as each major element is consumed in the centre, progressively heavier elements ignite, temporarily halting collapse. If the core of the star is not too massive (less than approximately 1.4 solar masses, taking into account mass loss that has occurred by this time), it may then form a white dwarf (possibly surrounded by a planetary nebula).

The most massive stars may be completely destroyed by a supernova with an energy greatly exceeding its gravitational binding energy. This rare event, caused by pair-instability, leaves behind no black hole remnant.

After a star has burned out its fuel supply, its remnants can take one of three forms, depending on the mass during its lifetime.

White and black dwarfs

For a star of 1 solar mass, the resulting white dwarf is of about 0.6 solar mass, compressed into approximately the volume of the Earth. White dwarfs are stable because the inward pull of gravity is balanced by the degeneracy pressure of the star's electrons. White dwarfs of higher mass have a smaller volume. With no fuel left to burn, the star radiates its remaining heat into space for billions of years.

In the end, all that remains is a cold dark mass sometimes called a black dwarf. However, the universe is not old enough for any black dwarf stars to exist yet.

If the white dwarf's mass increases above the Chandrasekhar limit, which is 1.4 solar masses then electron degeneracy pressure fails due to electron capture and the star collapses into a neutron star or a supernova marking the death of a massive star.

Neutron stars

When a stellar core collapses, the pressure causes electron capture, thus converting the great majority of the protons into neutrons. The electromagnetic forces keeping separate nuclei apart are gone (proportionally, if nuclei were the size of dust mites, atoms would be as large as football stadiums), and most of the core of the star becomes a dense ball of contiguous neutrons (in some ways like a giant atomic nucleus), with a thin overlying layer of degenerate matter.

These stars, known as neutron stars, are extremely small--on the order of radius 10 km, no bigger than the size of a large city--and are phenomenally dense. Their period of revolution shortens dramatically as the star shrinks (due to conservation of angular momentum); some spin at over 600 revolutions per second. When these rapidly rotating stars' magnetic poles are aligned with the Earth, we detect a pulse of radiation each revolution. Such neutron stars are called pulsars, and were the first neutron stars to be discovered.

Artist’s impression about the evolution of a star