SUPERNOVAE, TYPE II, THEORY AND INTERPRETATION STANFORD E. WOOSLEY A supernova (SN) of type II is the observable event corresponding to the death of a massive star. The class is defined by spectroscopic evidence for hydrogen in the ejecta, indicating that the star that explodes is one that has not lost its envelope. Although it is possible, in principle, for an intermediate mass star of roughly 5-8 M* (solar masses) to ignite carbon burning while still maintaining a partially intact envelope and thus to provide a type II event with a thermonuclear power source, just as in type I events, another mechanism is generally deemed responsible for type II supernovae: gravitational collapse. This paradigm received welcome confirmation recently when a burst of neutrinos was detected at the onset of SN 1987A. Such a neutrino burst could only have been generated by the collapse of the iron core of a massive evolved star to a neutron star. From photographs of the region taken before the explosion, it has been determined that the progenitor of SN 1987A had a mass on the main sequence near 20 M*. In theory, the lightest star that can experience core collapse to a neutron star and thus power a type II supernova is near 8 M*, a value that is somewhat sensitive to how convection is treated in the evolutionary calculation. The highest mass star that makes a type II supernova is determined by the efficiency of the still poorly understood supernova mechanism and by the possibility that stars heavier than about 40 M* may lose their entire hydrogen envelope either to a pulsational instability or to a radiatively driven wind. If such stars still were to succeed in exploding, they would be designated type Ib. THE LIVES OF MASSIVE STARS The life of a massive star is comparatively brief because such stars are very luminous and are profligate spenders of their nuclear energy reserves. A star of 10 M* will live for about 30 million years, a star of 20 M* for about 10 million years, and a 40 M* star, only 5 million years. Most of the time is spent burning hydrogen to helium. The remainder is spent in five other burning stages characterized by rapidly decreasing time scales. First helium burns to carbon and oxygen, then carbon to neon and magnesium, neon to (more) oxygen and magnesium, oxygen and magnesium to silicon and sulfur, and finally silicon and sulfur burn to elements of the iron group. Along the way, traces of many other elements are made, for example, sodium and aluminum in carbon burning, phosphorus in neon burning, and chlorine and potassium in oxygen burning. Each of these burning stages must surmount increasing charge barriers to nuclear fusion and so each stage occurs at a higher temperature than the previous one. Carbon burns at about one billion degrees Kelvin (*** K), but silicon burning requires 3.5x10* K. At temperatures this great, energy escapes from the core not only by the usual means of radiative diffusion and convection, but by neutrinos. The neutrinos are generated by the annihilation of electrons and positrons, themselves produced by copious gamma rays in the core. Because the neutrino losses scale as about the ninth power of the temperature and because nuclear resources are very limited, as fusion produces progressively heavier elements the time scales become shorter. Carbon may take 1000 years or more to burn, whereas oxygen burns in about one year and silicon burning takes only one week. For stars at the lower end (8-10 M*) of the type II supernova mass range, this standard scenario and the time scales are altered somewhat after carbon burning because the electrons become degenerate and provide an additional source of pressure in the core. In the end, however, the cores of all stars in the mass range we are considering collapse to neutron stars. At the end of silicon burning the (typical) massive star consists of an iron core of about 1.4 M* surrounded by layers consisting of the ashes of previous burning stages and a low-density hydrogen envelope (Fig. 1). The fraction consisting of helium and heavier elements ranges from roughly one-fourth (at 10 M*) to one-half (at 50 M*) of the original mass of the star. Smaller still is the fraction consisting of carbon and heavier elements. Once iron has been formed in the center, no further energy can be released by nuclear fusion. Gravitational contraction raises the temperature and density but not enough to provide the pressure needed to balance gravity. As the temperature of the iron core grows to 10x10** and more and the density to **********,electrons are squeezed into the nuclei, leading to heavier neutron-rich isotopes. Concurrently, high-energy radiation begins to tear nuclei apart into ** particles. Both processes rob the core of energy and pressure and although the pressure never actually decreases, it becomes weaker compared to gravity. The collapse of the core accelerates, eventually reaching speeds of up to 70,000 km s***. As the central density reaches and then exceeds the density of the atomic nucleus (****************) new forces and pressures come into play as the repulsive component of the strong, or nuclear, force brings an abrupt halt to the collapse. Roughly half of the core, that part in sonic communication, halts as a unit. The other half runs into this inner core at supersonic speed and bounces. A shock wave is born. A rebound of the compressed inner core as well as the energy from the reflecting material itself gives energy to this shock and it moves outwards. If enough energy is provided to the shock, it can exit the core, which is now making the transition to a neutron star, while retaining enough outward momentum and energy to eject the rest of the star with a kinetic energy near **** erg. THE EXPLOSION MECHANISM Unfortunately, most current calculations show that the shock loses so much energy to dissipative processes (neutrino losses and the photodissociation of bound nuclei into neutrons and protons) on the way out that by the time it reaches the edge of the core it has lost all outward kinetic energy. Were nothing else to intervene, there would be no supernova, just the relentless growth of the collapsed core by accretion to a state where not even the strong force could prevail against gravity. A black hole would be formed. However, the collapsed core, although very dense and neutron-rich, is not yet a neutron star. The binding energy of a cold neutron star (about 20% of its rest mass, or ****** erg) must still be radiated away over the next few seconds as neutrinos. If just a tiny fraction, only a few tenths of one percent of this neutrino energy, is deposited at the outer edge of the core (just behind where it has temporarily stalled and material is accreting), a powerful explosion still may develop. Detailed calculations by one research group have demonstrated this occurrence. Heating from neutrino energy deposition causes expansion, blowing a large bubble filled with radiation and electron-positron pairs. Expansion of this bubble causes the shock to move outward again with enough energy to eject all of the star external to the core with high velocity. The success or failure of this "delayed mechanism," so called because it takes roughly one second to develop in contrast to the 20 ms characterizing the shock crossing time of the core, depends upon a variety of nonthermal microscopic physical processes (neutrino scattering on electrons and nuclei, neutrino capture on neutrons and protons, and neutrino-neutrino annihilation, to name a few) and on macroscopic processes such as convection (and rotation and magnetic fields?), all coupled in a situation that can only be studied using sophisticated numerical codes. Not too surprisingly, no consensus has yet emerged as to the general validity of this mechanism. Almost certainly the explosion of some stars, if not all type II supernovae, occurs via neutrino energy transport, but there may well be additional physics that has yet to be modeled correctly in the codes. Whatever the situation on the computer, the existence of pulsars and supernovae assures us that some explosion mechanism works in nature. In order that material not fall back onto the neutron star and turn it into a black hole, a strong shock wave must somehow find its way into the heavy-element layers surrounding the collapsed core and expel them with high velocity. As this shock transits the shells of silicon and oxygen just outside the core, the high temperature it produces leads to a frenzy of nuclear fusion, producing a number of heavy elements from silicon through zinc. One of the most abundant of these is the radioactive nucleus **Ni, the most tightly bound of all nuclei having equal numbers of neutrons and protons (28). This nucleus decays with a half-life of 6.1 days to **Co, which in turn decays to **Fe with a half-life of 77.2 days. It is now believed that most iron in nature has been produced explosively as **Ni in both type I and type II supernovae and not as stable **Fe. Theory also suggests that from a few thousandths (10-M* supernova) to a few tenths (40-M*) of a solar mass of **Ni are produced in the explosion of each type II supernova. Analysis of SN 1987A has shown that 0.07 M* of **Ni was produced in that explosion. The decay of **Co is also a very important energy source to the supernova at late times (discussed below). THE SUPERNOVA APPEARS The shock wave moves on through the rest of the star, exiting the helium core in about one minute and the hydrogen envelope in from two hours to one day (depending upon whether the supernova occurred in a compact blue star as did SN 1987A or in a red supergiant as have most other type II supernovae). As the shock exits, the surface of the star is heated to temperatures near one-half million degrees and is accelerated to velocities of up to 30,000 km s**. It is at this point that the supernova first becomes visible to the (nonneutrino) astronomer. Early on, the emission is chiefly in the ultraviolet, but as the surface expands and cools the spectrum shifts into the optical. Within about one week the temperature near the surface of the supernova declines to near 6000 K and hydrogen begins to recombine. This recombination removes electrons that were the chief source of opacity and releases the energy that had been deposited in the hydrogen envelope by the shock wave. Recombination propagates as a front that, although carried outward in space by the expansion of the envelope, eventually moves inward in mass until the entire envelope has recombined. During this period, which may last from one week to three months depending upon the mass and radius of the envelope at the time the star explodes, the spectrum is similar to that of a blackbody having temperature 6000 K. Thus most of the emission is in optical wavelengths. Typically, a total of about 10** erg, or roughly 1% of the kinetic energy and 0.01% of the total neutrino energy, is radiated during this phase although the actual amount varies greatly from supernova to supernova. Once the envelope has recombined, the luminosity declines precipitously until a new energy source is found. That source is usually radioactivity: specifically, the energetic photons released as **Co decays to **Fe. Initially, all the gamma rays deposit their energy by scattering with electrons deep within the supernova. Then most of the energy still comes out at optical wavelengths. As time passes, however, an increasing fraction of the gamma rays (and x-rays from partially thermalized gamma rays) escape, although enough still deposit energy to keep the supernova shining brightly. X-rays and gamma rays having energy characteristic of **Co decay have been detected from SN 1987A. For a period of several years the luminosity of the supernova tracks the exponential decay of this radioactive isotope corrected for the partial escape of the gamma rays. At such late times other energy sources may also contribute to the light. These include the decay of radioactive isotopes besides **Co (especially **Co and **Ti), energy input by a pulsar, or energy from the supernova running into circumstellar material. Because the many factors upon which the light curve depends (the mass and radius of the hydrogen envelope, the mass of **Ni synthesized in the explosion, the presence of a pulsar or circumstellar material, and the energy of the explosion itself) are likely to vary from star to star, it is not surprising that the emission of type II supernovae is far less regular than that of type I supernovae. Spectroscopically, the emission of a type II supernova is dominated by lines of calcium, oxygen, and hydrogen. Velocities from 2000 to 30,000 km s** are inferred, with most of the mass (including the heavy elements made in the explosion) being ejected at the lower velocities, typically less than 4000 km s**. As the heavy elements expand, they cool. Providing that the density remains sufficiently high, a portion of the ejecta eventually may condense to form both molecules and dust. Emission from both has been seen in the late-time spectrum of SN 1987A. The presence of dust, as well as a decrease in the average excitation energy of collisionally excited heavy elements, results in a shift of a major fraction of the emission to the infrared at times later than about 600 days. Again, SN 1987A has provided an important example. Additional Reading Arnett, W.D., Bahcall, J.N., Kirshner, R.P., and Woosley, S.E. (1989). Supernova 1987A. Ann. Rev. Astron. Ap. 27 629. Behthe, H.A. and Brown, G.E.(1985). How a supernova explodes. Scientific American 252 (No. 5) 60. Petschek, A., ed.(1989). Supernovae. D. Reidel, Dordrecht. A collection of essays on the supernova phenomenon. Woosley, S.E. and Phillips, M.M.(1988). Supernova 1987A! Science 240 750. Woosley, S.E. and Weaver, T.A.(1986). The physics of supernova explosions. Ann. Rev. Astron. Ap. 24 205. Woosley, S.E. and Weaver, T.A.(1989). The great supernova of 1987. Scientific American 261 (No. 2) 32.