Stars as Laboratories for Fundamental Physics - MPP Theory Group

Anomalous Stellar Energy Losses Bounded by Observations 37
2.1.7 Intermediate-Mass Stars
The evolutionary scenario described so far applies to stars with a mass
below 2−3 M ⊙ . For larger masses the core conditions evolve continuously
to the ignition of helium which is then a quiet process. For a
given chemical composition the transition between the two scenarios is
a sharp function of the total mass. Because red giants are very bright
they dominate the total luminosity of an old stellar population so that
this transition affects its integrated brightness in a discontinuous way.
Some authors have used the concept of an “RGB phase transition” to
describe this phenomenon (Sweigart, Greggio, and Renzini 1990, and
references therein; Bica et al. 1991).
Stars with masses of up to 6−8 M ⊙ are expected to end up as CO
white dwarfs. Their evolution on the AGB can involve many interesting
phenomena such as “thermal pulses”—for a key to the literature see
Iben and Renzini (1984). Intermediate-mass stars have thus far played
little role for the purposes of particle astrophysics, and so their evolution
does not warrant further elaboration in the present context.
2.1.8 Massive Stars and Type II Supernovae
The course of evolution for massive stars (M > 6−8 M ⊙ ) is qualitatively
different because they ignite carbon and oxygen in their core,
allowing them to evolve further. This is possible because even after
mass loss they are left with enough mass that their CO core grows
toward the Chandrasekhar limit. Near that point the density is high
enough to ignite carbon which causes heating and thus temporarily relieves
the electron degeneracy. Next, the ashes of carbon burning (Ne,
Mg, O, Si) form a degenerate core which ultimately ignites Ne burning,
and so forth. Ultimately, the star has produced a degenerate iron core,
surrounded by half a dozen “onion rings” of different burning shells.
The game is over when the iron core reaches its Chandrasekhar limit
because no more nuclear energy can be released by fusion. The temperature
is at 0.8×10 10 K = 0.7 MeV, the density at 3×10 9 g cm −3 , and
there are about Y e = 0.42 electrons per baryon. Further contraction
leads to a negative feedback on the pressure as photons begin to dissociate
iron, a process which consumes energy. Electrons are absorbed
and converted to neutrinos which escape, lowering the electron Fermi
momentum and thus the pressure. Therefore, the core becomes unstable
and collapses, a process which is intercepted only when nuclear
density (3×10 14 g cm −3 ) is reached where the equation of state stiffens.