The Milky Way and Beyond (20 page)

The superabundant elements in the S-type stars come from the slow neutron process. Moreover, the observation of technetium-99 is ample evidence that these processes are at work in stars today. Even so, some low-abundance atomic nuclei are proton-rich (i.e., neutron-deficient) and cannot be produced by either the
s
or the
r
process. Presumably, they have been created in relatively rare events—e.g., one in which a quantum of hard radiation, a gamma-ray photon, causes a neutron to be ejected.

In addition, no known nuclear process is capable of producing lithium, beryllium, and boron in stellar interiors. These lightweight nuclei are probably produced by the breakdown, or spallation, of heavier elements, such as iron and magnesium, by high-energy particles in stellar atmospheres or in the early stages of star formation. Apparently, these high-energy particles, called cosmic rays, originate by means of electromagnetic disturbances in the neighbourhood of starspots and stellar flares, and they also arise from supernova explosions themselves. Some of these light-element nuclei also might be produced by cosmic rays shattering atoms of carbon, nitrogen, oxygen, and other elements in the interstellar medium.

Finally, the peculiar A-type stars comprise a class of cosmic objects with strange elemental abundance anomalies. These might arise from mechanical effects—for example, selective radiation pressure or photospheric diffusion and element separation—rather than from nuclear effects. Some stars show enhanced silicon, others enhanced lanthanides. The so-called manganese stars show great overabundances of manganese and gallium, usually accompanied by an excess of mercury. The latter stars exhibit weak helium lines, low rotational velocities, and excess amounts of gallium, strontium, yttrium, mercury, and platinum, as well as absences of such elements as aluminum and nickel. When these types of stars are found in binaries, the two members often display differing chemical compositions. It is most difficult to envision plausible nuclear events that can account for the peculiarities of these abundances, particularly the strange isotope ratios of mercury.

The final stages in the evolution of a star depend on its mass and angular momentum and whether it is a member of a close binary.

The faint white dwarf stars represent the endpoint of the evolution of intermediate- and low-mass stars. White dwarf stars, so called because of the white colour of the first few that were discovered, are characterized by a low luminosity, a mass on the order of that of the Sun, and a radius comparable to that of Earth. Because of their large mass and small dimensions, such stars are dense and compact objects with average densities approaching 1,000,000 times that of water.

Unlike most other stars that are supported against their own gravitation by normal gas pressure, white dwarf stars are supported by the degeneracy pressure of the electron gas in their interior. Degeneracy pressure is the increased resistance exerted by electrons composing the gas, as a result of stellar contraction. The application of the so-called Fermi-Dirac statistics and of special relativity to the study of the equilibrium structure of white dwarf stars leads to the existence of a mass-radius relationship through which a unique radius is assigned to a white dwarf of a given mass; the larger the mass, the smaller the radius. Furthermore, the existence of a limiting mass is predicted, above which no stable white dwarf star can exist. This limiting mass, known as the Chandrasekhar limit, is on the order of 1.4 solar masses. Both predictions are in excellent agreement with observations of white dwarf stars.

The central region of a typical white dwarf star is composed of a mixture of carbon and oxygen. Surrounding this core is a thin envelope of helium and, in most cases, an even thinner layer of hydrogen. A very few white dwarf stars are surrounded by a thin carbon envelope. Only the outermost stellar layers are accessible to astronomical observations.

White dwarfs evolve from stars with an initial mass of up to three or four solar masses or even possibly higher. After quiescent phases of hydrogen and helium burning in its core—separated by a first red-giant phase—the star becomes a red giant for a second time. Near the end of this second red-giant phase, the star loses its extended envelope in a catastrophic event, leaving behind a dense, hot, and luminous core surrounded by a glowing spherical shell. This is the planetary-nebula phase. During the entire course of its evolution, which typically takes several billion years, the star will lose a major fraction of its original mass through stellar winds in the giant phases and through its ejected envelope. The hot planetary-nebula nucleus left behind has a mass of 0.5–1.0 solar mass and will eventually cool down to become a white dwarf.

White dwarfs have exhausted all their nuclear fuel and so have no residual nuclear energy sources. Their compact structure also prevents further gravitational contraction. The energy radiated away into the interstellar medium is thus provided by the residual thermal energy of the nondegenerate ions composing its core. That energy slowly diffuses outward
through the insulating stellar envelope, and the white dwarf slowly cools down. Following the complete exhaustion of this reservoir of thermal energy, a process that takes several additional billion years, the white dwarf stops radiating and has by then reached the final stage of its evolution and becomes a cold and inert stellar remnant. Such an object is sometimes called a black dwarf.

White dwarf stars are occasionally found in binary systems, as is the case for the white dwarf companion to the brightest star in the night sky, Sirius. Aside from playing an essential role in Type Ia supernovae, they are also behind the outbursts of novae and of other cataclysmic variable stars.

Novae are a class of exploding stars whose luminosity temporarily increases from several thousand to as much as 100,000 times its normal level. A nova reaches maximum luminosity within hours after its outburst and may shine intensely for several days or occasionally for a few weeks, after which it slowly returns to its former level of luminosity. Stars that become novae are nearly always too faint before eruption to be seen with the unaided eye. Their sudden increase in luminosity, however, is sometimes great enough to make them readily visible in the nighttime sky. To observers, such objects may appear to be new stars; hence the name nova from the Latin word for “new.”

Most novae are thought to occur in double-star systems in which members revolve closely around each other. Both members of such a system, commonly called a close binary star, are aged: one is a red giant and the other a white dwarf. In certain cases, the red giant expands into the gravitational domain of its companion. The gravitational field of the white dwarf is so strong that hydrogen-rich matter from the outer atmosphere of the red giant is pulled onto the smaller star. When a sizable quantity of this material accumulates on the surface of the white dwarf, a nuclear explosion occurs there, causing the ejection of hot surface gases on the order of 1/10,000 the amount of material in the Sun. According to the prevailing theory, the white dwarf settles down after the explosion; however, the flow of hydrogen-rich material resumes immediately, and the whole process that produced the outburst repeats itself, resulting in another explosion about 1,000 to 10,000 years later. Recent research, however, suggests that such outbursts may recur at much longer intervals—every 100,000 years or so. It is explained that a nova eruption separates the members of the binary system, interrupting the transfer of matter until the two stars move close together again after a considerable length of time.

When the mass of a star's remnant core lies between 1.4 and about 2 solar masses, it apparently becomes a neutron star with a density more than a million times greater than even that of a white dwarf. These extremely dense, compact stars are thought to be composed primarily of neutrons. Neutron stars are typically about 20 km (12 miles) in diameter. Their masses range between 1.18 and 1.44 times that of the Sun, but most are 1.35 times that of the Sun. Thus, their mean densities are extremely high—about 10
¹⁴
times that of water. This approximates the density inside the atomic nucleus, and in some ways a neutron star can be conceived of as a gigantic nucleus.

Geminga pulsar, imaged in X-ray wavelengths by the Earth-orbiting XMM-Newton X-ray Observatory. The pair of bright X-ray “tails” outline the edges of a cone-shaped shock wave produced by the pulsar as it moves through space nearly perpendicular to the line of sight (from lower right to upper left in the image)
. European Space Agency

It is not known definitively what is at the centre of the star, where the pressure is greatest; theories include hyperons, kaons, pions, and strange quark matter. The intermediate layers are mostly neutrons and are probably in a “superfluid” state. The outer 1 km (0.6 mile) is solid, in spite of the high temperatures, which can be as high as 1,000,000 K. The surface of this solid layer, where the pressure is lowest, is composed of an extremely dense form of iron.

Another important characteristic of neutron stars is the presence of very strong magnetic fields, upwards of 10
¹²
Gauss (Earth's magnetic field is 0.5 Gauss), which causes the surface iron to be polymerized in the form of long chains of iron atoms. The individual atoms become compressed and elongated in the direction of the magnetic field and can bind together end-to-end. Below the surface, the pressure becomes much too high for individual atoms to exist.

The discovery of pulsars provided the first evidence of the existence of neutron stars. Pulsars are neutron stars that emit pulses of radiation once per rotation. The radiation emitted is usually radio waves. However, some objects are known to give off short rhythmic bursts of visible light, X-rays, and gamma radiation as well, and others are “radio-quiet” and emit only at X- or gamma-ray wavelengths. The very short periods of, for example, the Crab (NP 0532) and Vela pulsars (33 and 83 milliseconds, respectively) rule out the possibility that they might be white dwarfs. The pulses result from electrodynamic phenomena generated by their rotation and their strong magnetic fields, as in a dynamo. In the case of radio pulsars, neutrons at the surface of the star decay into protons and electrons. As these charged particles are released from the surface, they enter the intense magnetic field (10
¹²
Gauss; Earth's magnetic field is 0.5 Gauss) that
surrounds the star and rotates along with it. Accelerated to speeds approaching that of light, the particles give off electromagnetic radiation by synchrotron emission. This radiation is released as intense radio beams from the pulsar's magnetic poles.

These magnetic poles do not coincide with the rotational poles, and so the rotation of the pulsar swings the radiation beams around. As the beams sweep regularly past Earth with each complete rotation, an evenly spaced series of pulses is detected by ground-based telescopes.

Antony Hewish and Jocelyn Bell, astronomers working at the University of Cambridge, first discovered pulsars in 1967 with the aid of a radio telescope specially designed to record very rapid fluctuations in radio sources. Subsequent searches have resulted in the detection of about 2,000 pulsars. A significant percentage of these objects are concentrated toward the plane of the Milky Way Galaxy, the enormous galactic system in which Earth is located.

Although all known pulsars exhibit similar behaviour, they show considerable variation in the length of their periods—i.e., the intervals between successive pulses. The period of the slowest pulsar so far observed is about 11.8 seconds in duration. The pulsar designated PSR J1939+2134 was the fastest-known for more than two decades. Discovered in 1982, it has a period of 0.00155 second, or 1.55 milliseconds, which means it is spinning 642 times per second. In 2006 an even faster one was reported; known as J1748−2446ad, it has a period of 1.396 milliseconds, which corresponds to a spin rate of 716 times per second. These spin rates are close to the theoretical limit for a pulsar because a neutron star rotating only about four times faster would fly apart as a result of “centrifugal force” at its equator, notwithstanding a gravitational pull so strong that the star's escape velocity is about half the speed of light.

These fast pulsars are known as millisecond pulsars. They form in supernovae like slower rotating pulsars; however, millisecond pulsars often occur in binary star systems. After the supernova, the neutron star accretes matter from its companion, causing the pulsar to spin faster.

Careful timing of radio pulsars shows that they are slowing down very gradually at a rate of typically a millionth of a second per year. The ratio of a pulsar's present period to the average slow-down rate gives some indication of its age. This so-called characteristic, or timing, age can be in close agreement with the actual age. For example, the Crab Pulsar, which was formed during a supernova explosion observed in 1054 CE, has a characteristic age of 1,240 years; however, pulsar J0205+6449, which was formed during a supernova in 1181 CE, has a characteristic age of 5,390 years.