High mass stars in the range >8 solar masses

7.8 High mass stars in the range >8 solar masses

Stars in this mass range have suffi cient mass overlying the core so that the tem- perature of the core can increase beyond that in less massive stars. This allows the capture of alpha particles to proceed further. Having made carbon and oxygen,

Introduction to Astronomy and Cosmology

it is then possible to build up the heavier elements having atomic numbers increasing by four – produced by the absorption of alpha particles. Hence, the

16 O fuses to 20 Ne, the 20 Ne fuses to 24 Mg and then 24 Mg fuses to 28 Si producing a core dominated by silicon.

For each successive reaction to take place the temperature has to increase as there is a greater potential barrier for the incoming alpha particle to tunnel through. In the melee, protons can react with these elements to form nuclei of

other atoms with intermediate atomic numbers, such as 19 Fl and 23 Na, though these will be less common. A shell like structure results, with layers of the star containing differing elements, the heaviest nearest to the centre.

When the temperatures reach the order of 3 ⫻ 10 9 K, silicon can be trans- formed though a series of reactions passing through 32 S, 36 A and continuing up to

56 N. The silicon burning produces a core composed mostly of iron (the majority) and nickel. Iron and its close neighbours in the atomic table have the most stable

nuclei, and any further reactions to build up heavier nuclei are endothermic (they would absorb energy rather than provide it) so this is where nuclear fusion has to stop. The star is then said to have an iron core. This core is surrounded by shells in which the lighter elements are still burning giving an interior like that shown in Figure 7.7.

Figure 7.7 The ‘onion-like’ shells of fusion burning during the later stages in the evolution of a giant star. Image: Wikipedia Commons.

Stellar Evolution – The Life and Death of Stars

Figure 7.8 The binding energy curve.

The energy released by each stage of burning is reduced and, as the ele- ment involved lies nearer to the peak of the binding energy curve (Figure 7.8), progressively less energy is released per gram of fuel. As a result, the time spent carrying out each successive reaction becomes shorter: a star of mass 20 times that of our Sun will spend about 10 million years on the main sequence burning hydrogen to helium, then spend about 1 million years burning helium to carbon and about 300 years burning carbon to oxygen. The oxygen burning takes around 200 days and silicon burning is completed in just 2 days!

Once the core reaches its iron state, things progress very rapidly. At the tempera- tures that exist in the core (of order 8 ⫻ 10 9 K for a 15 solar mass star) the photons

have suffi cient energy to break up the heavy nuclei, a process known as photodis-

integration . An iron nucleus may produce 13 helium nuclei in the reaction:

56 Fe ⫹ γ → 13 4 He ⫹ 4n

These helium nuclei then break up to give protons and neutrons:

4 He ⫹ γ → 2 p ⫹ ⫹ 2n

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As energy is released when the heavy elements were produced, these inverse processes are highly endothermic (requiring energy to progress) and thus the temperature drops catastrophically. There is then not suffi cient pressure to support the core of the star which begins to collapse to form what is called a neutron star . In the forming neutron star, free electrons combine with the protons produced by the photodisintegration of helium to give neutrons, in the reaction:

p ⫹ ⫹e ⫺ →n⫹ν e

The electron neutrinos barely interact with the stellar material, so can immediately leave the star carrying away vast amounts of energy – the neu- trino luminosity of a 20 solar mass star exceeds it photon luminosity by seven orders of magnitude for a brief period of time! The outer parts of the core col- lapses at speeds up to 70 000 km s ⫺1 and, within about 1 s, the core, whose initial size was similar to the Earth, is compressed to a radius of about 40 km! This is so fast that the outer parts of the star, including the oxygen, carbon, and helium burning shells, are essentially left suspended in space and begin to infall towards the core.

The core collapse continues until the density of the inner core reaches about three times that of an atomic nucleus (∼8 ⫻ 10 14 g cm ⫺3 ). At this density, the strong nuclear force, which in nuclei is attractive, becomes repulsive – an effect caused by the operation of the Pauli exclusion principle to neutrons and termed neutron degeneracy pressure . As a result of this pressure, the core rebounds and a shock wave is propagated outwards into the infalling outer core of the star. As the material above is now so dense, not all the neutrinos escape immediately and give the shock front further energy which then continues to work its way out to the

surface of the star – there producing a peak luminosity of roughly 10 9 times that of our Sun. This is comparable with the total luminosity of the galaxy in which the star resides!