The evolution of a member of a close double-star system can be markedly affected by the presence of its companion. As the stars age, the more massive one swells up more quickly as it moves away from the main sequence. It becomes so large that its outer envelope falls under the gravitational influence of the smaller star. Matter is continuously fed from the more rapidly evolving star to the less massive one, which still remains on the main sequence. U Cephei is a classic example of such a system for which spectroscopic evidence shows streams of gas flowing from the more highly evolved star to the hotter companion, which is now the more massive of the two. Eventually, the latter will also leave the main sequence and become a giant star, only to lose its outer envelope to the companion, which by that time may have reached the white dwarf stage.

Novas appear to be binary stars that have evolved from contact binaries of the W Ursae Majoris type, which are pairs of stars apparently similar to the Sun in size but revolving around one another while almost touching. One member may have reached the white dwarf stage. Matter fed to it from its distended companion appears to produce instabilities that result in violent explosions or nova outbursts. The time interval between outbursts can range from a few score years to hundreds of thousands of years.

In ordinary novas the explosion seems to involve only the outer layers, as the star later returns to its former brightness; in supernovas the explosion is catastrophic. Normally, novas are small blue stars much fainter than the Sun, though very much hotter. When an outburst occurs, the star can brighten very rapidly, by 10 magnitudes or more in a few hours. Thereafter it fades; the rate of fading is connected with the brightness of the nova. The brightest novas, which reach absolute magnitudes of about −10, fade most rapidly, whereas a typical slow nova, which reaches an absolute magnitude of −5, can take 10 or 20 times as long to decline in brightness. This property, when calibrated as the absolute magnitude at maximum brightness versus the time taken to decline by two magnitudes, allows novas to be used as distance indicators for nearby galaxies. The changes in light are accompanied by pronounced spectroscopic changes that can be interpreted as arising from alterations in an ejected shell that dissipates slowly in space. In its earliest phases, the expanding shell is opaque. As its area grows, with a surface temperature near 7,000 K, the nova brightens rapidly. Then, near maximum light, the shell becomes transparent, and its total brightness plummets rapidly, causing the nova to dim.

The mass of the shell is thought to be rather small, about 10–100 times the mass of Earth. Only the outer layers of the star seem to be affected; the main mass settles down after the outburst into a state much as before until a new outburst occurs. The existence of repeating novas, such as the star T Coronae Borealis, suggests that perhaps all novas repeat at intervals ranging up to thousands or perhaps millions of years; and probably, the larger the explosion, the longer the interval. There is strong evidence that novas are components of close double stars and, in particular, that they have evolved from the most common kind of eclipsing binaries, those of the W Ursae Majoris type.

Stars of the SS Cygni type, also known as dwarf novas, undergo novalike outbursts but of a much smaller amplitude. The intervals between outbursts are a few months to a year. Such variables are close binaries. The development of this particular type may be possible only in close binary systems.

There are two major types of supernovas, designated Type I (or SNe I) and Type II (or SNe II). They can be distinguished by the fact that Type II supernovas have hydrogen features in their spectra, while Type I supernovas do not. Type II supernovas arise from the collapse of a single star more massive than about eight solar masses, resulting in either a neutron star or a black hole. There are three classes of Type I supernovas: Ia, Ib, and Ic. Type Ia supernovas are believed to originate in a binary system containing a white dwarf, rather like the case of ordinary novas. Unlike the latter, however, in which only the outer layers of the white dwarf seem to be affected, in a Type I supernova the white dwarf is probably completely destroyed; the details are not yet fully understood. Certainly, a supernova’s energy output is enormously greater than that of an ordinary nova. Type Ib and Ic supernovas are like Type II in that they are each the core collapse of a massive star. However, a Type II supernova retains its hydrogen envelope, whereas the Type Ib and Ic supernovas do not, thus leading to the different hydrogen features in their spectra. Type Ib retains a helium shell and so has a spectrum rich in helium lines; Type Ic does not retain the hydrogen or helium shell.

Empirical evidence indicates that in a Type Ia supernova the absolute magnitude at maximum light can be determined by a combination of data derived from the rate of dimming after maximum, the shape of the light curve, and certain colour measurements. A comparison of the absolute and apparent magnitudes of maximum light in turn allows the distance of the supernova to be found. This is a matter of great usefulness because Type Ia supernovas at maximum light are the most luminous “standard candles” available for determining distances to external galaxies and thus can be observed in more distant galaxies than any other kind of standard candle. In 1999, application of this technique led to the totally unexpected discovery that the expansion of the universe is accelerating rather than slowing down. This acceleration is caused by dark energy, a gravitationally repulsive force that is the dominant component (73 percent) of the universe’s mass-energy.

A peculiar explosive variable with no known counterpart is Eta Carinae (NGC 3372), which appears in telescopes on Earth as a fuzzy red “star” slightly less than two seconds of arc in diameter. Surrounding it is a shell of gas and dust shaped roughly like an hourglass divided by a thin disk. First observed as a star of about the fourth magnitude in 1677, it brightened irregularly, undergoing an outburst in 1843, when it became for a few years the second brightest star in the sky. Thereafter it slowly faded, becoming too faint for the unaided eye around the turn of the 20th century. The fading was due, at least in part, to obscuration by dust emitted in the earlier eruption. The star remained near seventh magnitude with irregular variations for most of the 20th century, but it began brightening again by one or two tenths of a magnitude per year in the mid-1990s. In 2005 astronomers found that Eta Carinae is, in fact, a binary star system with an orbital period of 5.52 years. Its A component has a temperature of about 15,000 K; its B component, about 35,000 K. Eta Carinae is considered to be one of a small class of stars known as luminous blue variables. Its luminosity has been estimated as five million times that of the Sun. Flaring events producing not only visible effects but also X-ray, ultraviolet, and radio-wave effects have been observed.

Probably all variable stars represent more or less ephemeral phases in the evolution of a star. Aside from catastrophic events of the kind that produce a supernova, some phases of stellar variability might be of such brief duration as to permit recognizable changes during an interval of 50–100 years. Other stages may require many thousands of years. For example, the period of Delta Cephei, the prototype star of the Cepheid variables, has barely changed by a detectable amount since its variability was discovered in 1784.