Astronomy/Stellar Evolution

For 2012, the Astronomy event will focus on stellar evolution and type Ia supernovae.

Birth
The life cycle differs between stars depending on their mass. Normal-mass stars begin in stellar nurseries, and some matter condenses to create a protostar. This gains more mass until fusion (H -> He) begins, when it becomes a main-sequence star.

Protostar
When a star is in free-fall collapse, it is a protostar. Its evolution revolves around the point that its energy comes from gravitational contraction. Due to its larger radius, it is more luminous than it will be on the main sequence. Its low temperature allows for low opacity; due to this low opacity, its main form of energy transport is convection. Protostars gradually shrink and gain material as their central temperature rises, then drops. As it contracts, it moves down on the H-R diagram. Its temperature doesn't change much at first, but when it rises, opacity decreases, allowing radiation to take over. When the star loses much energy to radiation, it moves to the left on the H-R diagram.

Very low mass stars with masses less than 0.08 solar masses cannot reach the main sequence. They may become brown dwarfs or, more rarely, planets.

T Tauri Star
T Tauri stars are pre-main-sequence variable stars with spectral classes from F to M. They have many emission lines in their spectra, indicating their strong stellar winds. They are easy to identify and can be used as traces of solar-mass star formation regions. T Tauri stars almost always appear within dark dust clouds.

Main Sequence
Stars spend the majority of their lives (about 80 percent) at this stage. The main sequence lifetime of a solar-mass star is approximately 10 billion years. How long a star remains on the main sequence depends on its mass.

A star becomes a main sequence star when it is obtaining all its radiated energy from nuclear fusion of hydrogen into helium. The pressure that maintains hydrostatic equilibrium with gravity is exerted by collisions within the star. These collisions will eventually excite electrons in atoms, which then emit light waves. Some of these waves will escape the star; this lost energy comes at the price of lost random motion of atoms, which caused collisions in the first place. If nothing could replace this energy, the star would contract because pressure would decrease and gravity would take over. This is why thermonuclear reactions in the form of hydrogen burning must occur.

Lower-mass stars burn hydrogen using the proton-proton chain. In this reaction, four hydrogen atoms combine to form a helium atom. The rate of this reaction is equal to $$T^4$$, where T is temperature.

Such stars may also produce some energy by means of the CNO cycle, discussed in further depth under massive stars.

Maturity
When a main sequence star exhausts its core hydrogen, thermonuclear reactions cease. Gravity takes over and contracts the core. This heats the layer of hydrogen, so it can burn in a shell around the core in shell burning. This burning heats the surrounding areas, making them expand. Their temperatures decrease as their radii increase. The decrease in temperature increases opacity, resulting in convection taking over radiation again. The star's luminosity increases significantly; it moves up the red giant branch on the H-R diagram, a region of higher luminosities and lower surface temperatures. 50 million years after leaving the main sequence, a red giant has a surface temperature of 3000 K and a radius of nearly 100 solar orbits. Gravity compresses the core, about 10 earth radii in diameter, so much that electrons form a degenerate gas (that is, it has such high density that it does not obey the ideal gas law; it resists pressure). This allows the star to maintain hydrostatic equilibrium without thermonuclear reactions.

A helium core builds up inside the star, but there isn't enough heat/pressure to fuse the helium into heavier elements. However, the hydrogen "shell" around the helium core starts to fuse at a higher rate, causing the star to expand into a red giant and become more luminous. Then, as the star uses up its store of hydrogen, the outer layers of the star contract, finally achieving enough heat and pressure for the He in the core to fuse to carbon and oxygen. Helium burning occurs by the triple-alpha process. Because the gas in the core is degenerate, once it has been ignited, fusion spreads rapidly throughout the core. The temperature increases, increasing the rate of reactions, which then increases the temperature again in a cycle. This causes a helium flash when helium is suddenly ignited.

Eventually, the triple-alpha process converts the nuclei in the core to carbon. After a slight contraction to heat the remaining helium, helium burning continues in a shell in a manner like that of hydrogen burning. The burning shells make the star expand enough to reach the red giant phase again, which it left after ceasing helium burning. Electrons form a degenerate gas again; the star moves up the asymptotic giant branch on the H-R diagram.

Because the rate of the triple-alpha process is highly sensitive to temperature, the heat of the helium-burning shell makes the star unstable. The star will contract a little, increasing temperature, energy production, and pressure in the helium layer. The pressure increase overcompensates for gravity, so the star expands. This expansion then decreases the temperature, energy production, and pressure, so gravity contracts the star again. This continues in cycles known as thermal pulses.

Death
Normal-mass stars don't have enough mass to fuse carbon and oxygen into any heavier elements. During the thermal pulses, the star has a strong outflow of mass called a superwind. Once the star uses its entire store of hydrogen and helium, the outer layers of the star are ejected at high speed, potentially forming a planetary nebula. The remaining carbon/oxygen core heats these layers, exciting photons so the nebula glows due to fluorescence. The core then cools to become a white dwarf. If a white dwarf accumulates enough mass (perhaps gas from its partner in a binary system), it will explode in a Type Ia supernova. The "mass limit", so to speak, of a white dwarf is called the Chandrasekhar Limit and is equal to about 1.4 solar masses.

High-mass stars
Larger stars are similar, except they begin with more mass and grow to supergiants. However, high-mass stars DO have enough mass to fuse carbon and oxygen into heavier elements, each step of which temporarily creates enough outwards pressure to keep the star from collapsing under its own mass. Fusion continues all the way to iron in a process known as nucleosynthesis. Any element heavier than iron releases energy through fission instead of fusion. At the end of their lifetime, they can explode in a massive explosion known as a supernova and/or collapse into a neutron star or a black hole.

As main sequence stars, high-mass stars produce energy through the CNO cycle. This form of hydrogen burning uses carbon, nitrogen, and oxygen as catalysts for the production of helium.

Low-mass stars
Smaller mass stars (red dwarfs) don't become giants. Due to a main sequence lifespan longer than the age of the universe, no evolved red dwarf has been observed. Some current models predict the red dwarf increasing in surface temperature while maintaining a constant radius, transforming them into blue dwarfs. Upon the termination of nuclear fusion, the blue dwarf will cool into a white dwarf and eventually cool into a black dwarf.

In general, the mass of a star is inversely proportional to its lifespan - smaller stars (red dwarfs) live much longer than our own Sun (an "average" star), which in turn has a much longer lifespan than massive stars like Vega.

Stellar Evolution and the H-R Diagram
For an introduction to the basics of H-R Diagrams, please see the Astronomy main page

The H-R Diagram is very important with regards to stellar evolution because it can be used to identify the life cycle of a star, as well as characteristics of the stars in a star cluster.