Our simulation begins with stars plotted on the main sequence. On the HR diagram, the main sequence is an almost straight line from the upper left to the lower right region of the diagram. We do not concern ourselves with showing any of the very short-lived proto-stellar evolution; stars begin on the main sequence where they spend most of their lives. It is at this time in a star's life that it is in hydrostatic equilibrium and is turning hydrogen into helium and energy in its core in a process known as thermonuclear fusion. The star is very stable once it reaches this point and will exist on the main sequence until the core hydrogen fusion ceases. During this stable hydrogen fusion process, a star's properties will evolve very little. A low-mass star will increase in luminosity as well as temperature as it fuses its hydrogen, causing it to move slightly up along the main sequence (to the upper left). However, a high-mass star will increase in luminosity but decrease in temperature, causing it to slowly move off the main sequence to the right on the HR diagram (Hurley et al. 2000). There is also a correlation between a star's mass and its initial location on the main sequence. The higher a star is on the main sequence (i.e. the higher its luminosity and temperature), the higher its mass will be.
After most of the hydrogen in a star's core has been fused to helium, fusion stops, and the new helium core begins to contract due to its self-gravity. Hydrogen fusion still continues in a thin layer above the core. The star begins to evolve off of the main sequence at this point. As the helium core contracts, the layers above the core fuse hydrogen at a more rapid rate, because of increased temperature. This increase in hydrogen-shell burning causes the outer layers of the star to expand and cool and thus redden. At this point the star is becoming a red giant star (Chaisson and McMillan). As this occurs, the star moves to the right of the main sequence on the HR diagram. As the star leaves the main sequence, and just before it becomes a red giant, it will enter the Hertzsprung gap or subgiant branch. This occurs very rapidly, usually in less than a million years, and at an almost constant luminosity. Thus, it is difficult to actually observe stars in this phase of evolution (Hurley et al. 2000).
After the star leaves the short-lived subgiant branch, it enters the giant branch as the core continues to collapse and the outer layers continue to expand. Eventually, the helium core of the star will reach a critical temperature and pressure at which the helium will begin to fuse into carbon and oxygen. For a star of mass less than about 2 solar masses, the star will have a degenerate helium core at the time of helium fusion and will begin this new fusion process rather violently in what is known as the helium flash. The helium flash can be seen on the HR diagram as a spike in the movement of the lower mass stars' evolutionary tracks. Immediately after the helium flash, a star rapidly moves down the HR diagram (towards less luminosity) and then back towards the main sequence in a horizontal fashion (towards higher temperatures) into a period of stable helium fusion known as the horizontal branch. This transition onto the horizontal branch is more gradual for stars of mass greater than about 2 solar masses. For stars of mass greater than about 12 solar masses, it is possible for the core temperatures of these stars to rise quickly enough that helium fusion begins with almost no giant branch phase of evolution. This can be seen on the HR diagram as stars of this mass move horizontally away from the main sequence without significantly increasing luminosity (Chaisson and McMillan, Hurley et al. 2000).
The new helium fusion process does not last very long compared to the original hydrogen fusion process that occurred on the main sequence. The higher temperature in a star's core present at the time of helium fusion causes the rate of helium burning to be higher than that of the previous hydrogen core burning (Chaisson and McMillan). The luminosity is similar to what it is on the main sequence, even though helium fusion only releases 1/10 of the energy per unit mass that hydrogen fusion does. After all of the helium has been fused in the core of a star, the star undergoes changes similar to the changes it went through after the main sequence phase. The now carbon/oxygen core begins to contract under its own gravity as the outer layers of the star again expand and cool. This phase of stellar evolution is known as the asymptotic-giant branch and can be seen on the HR diagram as the path of the star as it moves away from the horizontal branch towards a region of lower temperature and in the case of lower mass stars, higher luminosity. The star is now becoming a red supergiant (Chaisson and McMillan, Hurley et al. 2000).
What happens next to the star depends critically on its mass. If the star's mass is less than about 8 and greater than about 0.3 solar masses, the star ends its life as a carbon/oxygen white dwarf (Hurley et al. 2000). As the carbon/oxygen core contracts, the outer layers become more and more unstable, eventually expanding off of the star completely and creating what is known as a planetary nebula. The remaining carbon/oxygen core contracts to the point at which electron degeneracy pressure counteracts the inward pull of gravity. This degeneracy pressure is not any kind of gas pressure and is independent of temperature, and only occurs at very high density (>106g/cm3). It is at this point that the core of the now dead star has become a white dwarf, an extremely dense star that has a high temperature but a low luminosity. The white dwarf region is at the lower left of the HR diagram, below the main sequence (Chaisson and McMillan, Hurley et al. 2000).
A star of mass less than about 0.3 solar masses never reaches high enough central temperatures to fuse helium into carbon and oxygen. However, after a long time, the star's outer layers will dissipate into a planetary nebula because of the low gravitational pull, revealing a mainly helium core that is known as a helium white dwarf (Chaisson and McMillan).
At the other extreme, a star of mass greater than about 8 solar masses will have a core too massive to be supported by electron degeneracy pressure after the carbon/oxygen core contracts. The star is massive enough that it goes through several more fusion processes (carbon and oxygen can fuse to neon, etc.). Each new fusion process gives less energy and lasts for less time, until the core has become primarily iron. As the star goes through these different fusion processes, it moves on the HR diagram somewhat horizontally away from and to the right of the main sequence. Once the core becomes iron, no fusion process can counteract the gravitational pull of the core on itself. Fusion of iron is an endothermic process, absorbing energy instead of releasing it, making it impossible for iron to be fused as an energy source for the star. Thus, the iron core collapses due to its own gravity, until the central densities reach 1014 g/cm3, near the densities of atomic nuclei. At these densities, a neutron degeneracy pressure (stronger than the electron degeneracy pressure) prevents further collapse. The inner core is now a proto neutron star, and the neutron degeneracy pressure rebounds the outer material. This rebound sends a shockwave up through the rest of the star, exploding the outer layers of the core as well as the outer layers of the star itself in a violent event known as a type II supernova (Hurley et al. 2000).
What happens to the remaining core after this explosion depends on the remaining core's mass. If the core's mass is less than about 3 solar masses, the core's neutron degeneracy pressure continues to counteract the inward pull of gravity, stabilizing it into a neutron star. Neutron degeneracy pressure occurs at much higher densities than electron degeneracy pressure, making neutron stars even denser than white dwarfs. Neutron stars are also dimmer than white dwarfs, putting them even lower on the HR diagram, but still in the same general region below the main sequence. Many neutron stars have strong remnant magnetic fields. Some of these neutron stars collect infalling matter at and radiate out of their magnetic poles, which we observe as pulsars. If the core's remaining mass is greater than about 3 solar masses, not even neutron degeneracy pressure will be able to withstand the inward pull of gravity, and the core will collapse into a black hole from which no electromagnetic radiation can escape (Chaisson and McMillan).
Stars of mass greater than about 15 solar masses go through one additional phase of evolution. These stars lose a significant amount of their initial mass to strong stellar winds. These stars may even lose their entire outer envelope during helium fusion or sometimes during the quick subgiant phase. If this occurs, the stars' helium burning cores will be revealed before they have become white dwarfs or neutron stars, making stars known as naked helium stars. They appear to the left of the main sequence at the very highest temperatures (>50,000 K).