A low-mass star—one with a mass of less than about eight solar masses—never becomes hot enough to burn carbon in its core. It ends its life as a carbon white dwarf. A high-mass star, however, can fuse not just hydrogen and helium but also carbon, oxygen, and even heavier elements as its inner core continues to contract and its central temperature continues to rise. 
(Sec. 20.4) The burning rate accelerates as the core evolves. Can anything stop this runaway process? Is there a stable "white-dwarf—like" state at the end of the evolution of a high-mass star? What is its ultimate fate? To answer these questions, we must look more carefully at fusion in massive stars.
Figure 21.5 is a cutaway diagram of the interior of a highly evolved star of large mass. Note the numerous layers where various nuclei burn. As the temperature increases with depth, the ash of each burning stage becomes the fuel for the next stage. At the relatively cool periphery of the core, hydrogen fuses into helium. In the intermediate layers, shells of helium, carbon, and oxygen burn to form heavier nuclei. Deeper down reside neon, magnesium, silicon, and other heavy nuclei, all produced by nuclear fusion in the layers overlying the core. The core itself is composed of iron. We will study the key reactions in this burning chain in more detail later in this chapter.
Figure 21.5 Cutaway diagram of the interior of a highly evolved star of mass greater than 8 solar masses. The interior resembles the layers of an onion, with shells of progressively heavier elements burning at smaller and smaller radii and at higher and higher temperatures.
As each element is burned to depletion at the center, the core contracts, heats up, and starts to fuse the ash of the previous burning stage. A new inner core forms, contracts again, heats again, and so on. Through each period of stability and instability, the star's central temperature increases, the nuclear reactions speed up, and the newly released energy supports the star for ever-shorter periods of time. For example, in round numbers, a star 20 times more massive than the Sun burns hydrogen for 10 million years, helium for 1 million years, carbon for 1000 years, oxygen for 1 year, and silicon for a week. Its iron core grows for less than a day.
Once the inner core begins to change into iron, our high-mass star is in trouble. Nuclear fusion involving iron does not produce energy—iron nuclei are so compact that energy cannot be extracted by combining them into heavier elements. In effect, iron plays the role of a fire extinguisher, damping the inferno in the stellar core. With the appearance of substantial quantities of iron, the central fires cease for the last time, and the star's internal support begins to dwindle. The star's foundation is destroyed, and its equilibrium is gone forever. Even though the temperature in the iron core has reached several billion kelvins by this stage, the enormous inward gravitational pull of matter ensures catastrophe in the very near future. Gravity overwhelms the pressure of the hot gas, and the star implodes, falling in on itself.
The core temperature rises to nearly 10 billion K. At these temperatures, individual photons, according to Wien's law, have tremendously high energies, enough to split iron into lighter nuclei and then to break those lighter nuclei apart until only protons and neutrons remain. (Sec. 3.4) This process is known as photodisintegration of the heavy elements in the core. In less than a second, the collapsing core undoes all the effects of nuclear fusion that occurred during the previous 10 million years! But to split iron and lighter nuclei into smaller pieces requires a lot of energy. After all, this splitting is just the opposite of the fusion reactions that generated the star's energy during earlier times. Photodisintegration absorbs some of the core's thermal energy—in other words, it cools the core and so reduces the pressure. As nuclei are destroyed, the core of the star becomes even less able to support itself against its own gravity. The collapse accelerates.
Now the core consists entirely of simple elementary particles—electrons, protons, neutrons, and photons—at enormously high densities, and it is still shrinking. As the core density continues to rise, the protons and electrons are crushed together, forming neutrons and neutrinos:
This process is sometimes called the neutronization of the core. Recall from our discussion in Chapter 16 that the neutrino is an extremely elusive particle that interacts hardly at all with matter. (Sec. 16.5) Even though the central density by this time may have reached 1012 kg/m3 or more, most of the neutrinos produced by neutronization pass through the core as if it weren't there. They escape into space, carrying away energy as they go.
The disappearance of the electrons and the escape of the neutrinos make matters even worse for the core's stability. There is now nothing to prevent it from collapsing all the way to the point at which the neutrons come into contact with one another, at the incredible density of about 1015 kg/m3. At this point the neutrons in the shrinking core play a role similar in many ways to that of the electrons in a white dwarf. When far apart they offer little resistance to compression, but when brought into contact, they produce enormous pressures that strongly oppose further gravitational collapse. This neutron degeneracy pressure, akin to the electron degeneracy pressure that operates in red giants and white dwarfs, finally begins to slow the collapse. (Sec. 20.2) By the time the collapse is actually halted, however, the core has overshot its point of equilibrium, and may reach densities as high as 1017 or 1018 kg/m3 before turning around and beginning to reexpand. Like a fast-moving ball hitting a brick wall, the core becomes compressed, stops, then rebounds—with a vengeance!
The events just described do not take long. Only about a second elapses from the start of the collapse to the "bounce" at nuclear densities. At that point the core rebounds. An enormously energetic shock wave sweeps through the star at high speed, blasting all the overlying layers—including the heavy elements outside the iron inner core—into space. Although the details of how the shock reaches the surface and destroys the star are still uncertain, the end result is not. The star explodes, in one of the most energetic events known in the universe (see Figure 21.6). For a period of a few days the exploding star may rival in brightness the entire galaxy in which it resides. This spectacular death rattle of a high-mass star is known as a core-collapse supernova.
Figure 21.6 A supernova called SN1987A (arrow) was exploding near this nebula (30 Doradus) at the moment the photographon the right was taken. The photograph on the left is the normal appearance of the star field. (Interlude 21-1.)
