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Nuclear processes

Radioactive decay processes Alpha decay Beta decay Cluster decay Double beta decay Double electron capture Electron capture Gamma radiation Internal conversion Isomeric transition Neutron emission Positron emission Proton emission Spontaneous fission Nucleosynthesis Nuclear fusion Proton-proton chain reaction CNO cycle Triple-alpha process Carbon burning process Neon burning process Oxygen burning process Silicon burning process Neutron capture The R-process The S-process Proton capture: The P-process The Rp-process Spallation

In astrophysics, silicon burning is a one-day-long sequence of nuclear fusion reactions that occur in “massive” stars with a minimum of about 8–11 solar masses. Silicon burning is an end-of-life process for stars that have run out of the fuels that power them for the long periods while they are the main sequence on the Hertzsprung-Russell diagram. Silicon burning begins when gravitational contraction raises the star’s core to a temperature of 2.7–3.5 billion kelvin (GK). The exact temperature depends on mass. When a star has completed the silicon-burning phase, it explodes in what is known as a Type II supernova.

Contents 1 Nuclear fusion sequence and the alpha process 2 Binding energy 3 See also 4 External links

Stars with normal mass (no greater than about three solar masses) run out of fuel after their hydrogen has been consumed and fused into helium. If the star has intermediate mass (greater than three solar masses but less than about eight) the star can “burn” (fuse) helium into carbon. These stars end their lives when their helium has been exhausted and they have a carbon core. High mass stars (>8–11 solar masses) are able to burn carbon because of the extraordinarily high gravitational potential energy bound in their mass. As a massive star contracts, its core heats up to 600 MK and carbon burning begins which creates new elements as follows:

carbon–12 → oxygen–16, neon–20, and magnesium–24

The chemical elements are defined by the number of protons in their nucleus. In the elements listed above, the suffix denotes a particular isotope (form of a chemical element having a different number of neutrons) in terms of its molar mass.

After a high-mass star has burned all its carbon, it contracts, gets hotter, and begins burning the oxygen, neon, and magnesium as follows:

oxygen–16, neon–20, and magnesium–24 → silicon–28 and sulfur–32 (a six-month-long process)

After high-mass stars have nothing but sulfur and silicon in their cores, they further contract until their cores reach in the range of 2.7–3.5 GK; silicon burning starts at this point. Silicon burning entails the alpha process which creates new elements by adding the equivalent of one helium nucleus (two protons plus two neutrons) per step in the following sequence:

silicon–28 → sulfur–32 → argon–36 → calcium–40 → titanium–44 → chromium–48 → iron–52 → nickel–56

The entire silicon-burning sequence lasts about one day and stops when nickel–56 has been produced. Nickel–56 (which has 28 protons) has a half-life of 6.02 days and decays via beta radiation (beta plus decay, which is the emission of a positron) to cobalt–56 (27 protons), which in turn has a half-life of 77.3 days as it decays to iron–56 (26 protons). However, only minutes are available for the nickel–56 to decay within the core of a massive star. At the end of the day-long silicon-burning sequence, the star can no longer convert mass into energy via nuclear fusion because a nucleus with of 56 nucleons has the lowest mass per nucleon (proton and neutron) of all the elements in the alpha process sequence; which is to say, the binding energy of nickel–56 has reached a maximum. Accordingly, fusing additional nucleons to nickel–56 would actually consume energy rather than release it. Within minutes, the star begins to contract. The potential energy of gravitational contraction heats the interior to 5 GK and this opposes and delays the contraction. However, since no additional heat energy can be generated via new fusion reactions, the contraction rapidly accelerates into a collapse lasting only a few seconds. The central portion of the star gets crushed into either a neutron star or, if the star is massive enough, a black hole. The outer layers of the star are blown off in an explosion known as a Type II supernova that lasts days to months. The supernova explosion releases large quantities of neutrons which synthesizes in about one second, roughly half the elements heavier than iron via a neutron-capture mechanism known as the r-process (where the “r” stands for rapid neutron capture).

The graph below shows the binding energy of various elements. Increasing values of binding energy can be thought of in two ways: 1) it is the energy required to remove a nucleon from a nucleus, and 2) it is the energy released when a nucleon is added to a nucleus. As can be seen, light elements such as hydrogen release large amounts of energy (a big increase in binding energy) as nucleons are added—the process of fusion. Conversely, heavy elements such as uranium release energy when nucleons are removed—the process of nuclear fission. Adding to a nucleus with 56 nucleons (such as iron–56 or nickel–56) requires energy rather than releasing any. Nickel–56 is the last fusion product produced in the core of a high-mass star.

• Stellar evolution
• Supernova nucleosynthesis
  • Neutron capture: p-process, r-process s-process

• Stellar Evolution: The Life and Death of Our Luminous Neighbors, by Arthur Holland and Mark Williams of the University of Michigan
• The Evolution and Death of Stars, by Ian Short
• Star, by World Book @ NASA
• Origin of Heavy Elements, by Tufts University
• Chapter 21: Stellar Explosions, by G. Hermann