Star Evolution

One of the main goals of the theory of stellar evolution is to understand, why stars cluster in certain regions of the H-R diagram, and how they evolve from one part to another. The H-R diagram is very useful in understanding the current stage of the evolution of a star. In star evolution the mass of a

Abbildung 2.2: Evolution of stars. The left picture shows the process of birth of stars and their evolution to the main-sequence. The place they settle on the main-sequence is determined by their initial masses. This process is also called as Hayashi contraction, and lasts millions of years. The picture on the right shows the process of dying of three different types of stars: massive stars, sun-like stars and dwarf stars. The evolution of time can be followed as indicated by arrows.

star plays a very important role, because stars with different masses follow different paths in the H-R diagram in their evolution, which can be seen in different phases that the star goes through.

While the star becomes stable, its position on the H-R diagram moves according to its mass, from the upper right corner of the diagram, which is the faint and cool stage of the star, to the upper left side, which is the hot and bright phase. Now, the star starts to evolve on the thermal time scale.

In this so-called pre-main-sequence phase, the star moves slowly from the upper left position on the HR-diagram to settle somewhere along the main-sequence stars depending on its mass. This phase is illustrated in Figure 2.2 (left) and it is also known as the Hayashi contraction phase of a protostar.

The time for a star to reach the main-sequence varies with its mass. A star with a mass of our sun (M) reaches the main-sequence in 3×10⁷ years. A star with a 0.5M come to this stage in 10⁸ years and another with a mass about 15M in 6×10⁵ years.

In themain-sequence phase, temperatures in the cores of the stars are so high that hydrogen starts to be converted to helium releasing 0.7% of the rest mass energy, which is the binding energy of helium. The primary output from these so-calledthermonuclearreactions are photons and a large number of other particles such as neutrinos. There are two types of thermonuclear

reactions, by which hydrogen can be converted into helium. The type of reaction is determined by the initial temperature of the core of the star.

Therefore, when the temperature of the star is less than about 2×10⁷ K, the p-p chain reactionis the primary energy source. If the temperature is greater than this value, the reaction cycle is known as the carbon-nitrogen-oxygen (CNO) cycle, which becomes a dominant process. In the p-p reaction the hydrogen is used as a catalyst, whereas in the CNO cycle the ¹²C is used as catalyst in the formation of helium (⁴He). The hydrogen burning phase is remarkably stable. For example, a solar mass star will live almost 10 billion years.

Thepost-main-sequenceevolution appears to be different for massive stars than for low-mass stars. When a low mass star (^<∼ 8M) exhausts the supply of hydrogen in its core, it contracts under gravity, heats up, and finally burns helium causing its luminosity to increase and move from the main-sequence to the giant branch (Figure 2.2 (right)). Eventually, helium burns completely and leaves the carbon core behind. Because the core can not withstand its own mass, it collapses under its own weight. At some stage the matter becomes so dense that the electron degeneracy pressure provides the balance against the weight. If the mass of the core is less than 1.2M, the star turns into a white dwarf and the core moves into the white dwarf branch of the H-R diagram.

The high-mass stars (∼^> 8M) continue to convert hydrogen into helium and helium into carbon and slowly move up to the giant branch in the H-R diagram. After the supply of helium in the core is depleted, lighter elements are fused to form heavier ones. Finally, iron is produced in the core, which is the most tightly bound element, but the production of iron continues in the surrounding layers. At some point gravitational pressure in the core ex-ceeds the electron degeneracy pressure and core collapse follows. Due to the core collapse, the temperature in the core rises and the photo-disintegration of iron into helium occurs. The newly formed helium atoms then further disintegrate into protons and neutrons. Protons in turn combine with am-bient electrons to form neutrons. Eventually, the neutron density increases and the neutron degeneracy pressure prevents further gravitational collapse.

Meanwhile, however, the outer layers continue falling inward and eventually rebound in a massive explosion. The result is a huge shock wave that moves radially out from the core expanding into interstellar space medium (ISM).

The surviving degenerate core is extremely dense, with typical mass of M ' 1.4M and radius R ' 15 km.

This type of explosion of a single star is called supernova explosion type II (The supernova explosion type I usually happens in binary systems.), the surviving core is referred to as the neutron star, and the shock-front is called

Also the vicinity of a neutron star which is highly magnetized, or a jet of an AGN are possible production sites of high-energy γ rays. In the followi-ng Sections, the possible production mechanisms of high-energy γ rays are briefly summarized.

2.2.1 Charged Particles in Strong Electric or Magnetic