A star is any massive, self-luminous celestial body of gas that shines by radiant energy generated within its interior. The universe contains trillions of stars, only a very small percentage of which are visible to the unaided eye.
The closest star to the solar system is Proxima Centauri, which is approximately 4.3 light-years from the Sun. The most distant stars lie in galaxies billions of light-years away. Stars may occur singly, as in the case of the Sun. More commonly, however, they exist in pairs, multiple systems with several members, or clusters consisting of numerous components. Furthermore, stars vary greatly in brightness, colour, temperature, mass, size, chemical composition, and age.
The observed brightness of a star is expressed by its apparent magnitude, with a decrease of one magnitude corresponding to an increase in brightness of 2.512 times (i.e., a star of the first magnitude is 2.512 times as bright as one of the second magnitude). Some stars are so bright as to be given negative magnitudes, as for example Sirius, with an apparent magnitude of -1.5. Stars visible to the unaided eye have values of six or lower. A star's apparent magnitude is not a direct indicator of its luminosity. Such intrinsic brightness is given in terms of absolute magnitude, which is the magnitude that a star would have if it were viewed at a standard distance of 10 parsecs (32.6 light-years).
Stars differ widely in colour, ranging from blue-white to red. Stellar colours can be measured by several methods. The most widely employed technique requires the use of a photoelectric photometer, which can measure stellar colours with the aid of filters.
The atmospheric, or surface, temperature of a star can be estimated by studying its spectrum, because temperature largely determines the types of absorption lines present. In nearly all stars, hydrogen is the most abundant element, but it is completely ionized in the outer layers of stars having a temperature in excess of 25,000 K and thus produces virtually no absorption lines. In cooler stars of about 10,000 K, many hydrogen atoms remain non-ionized, and a sufficient number are excited to higher energy levels to yield strong absorption lines (e.g., Balmer series). The spectra of stars with still lower surface temperatures manifest different absorption line patterns. The prevailing system of spectral classification divides stars into seven broad groups in order of decreasing temperature: O, B, A, F, G, K, and M. This O-M spectral sequence is also a colour sequence. For example, the O-type stars, which are the hottest of all stellar bodies, have an estimated surface temperature of more than 25,000 K and are intrinsically blue in colour. G-type stars, including the Sun, are relatively cool, with a temperature of 5,000-6,000 K, and are white to yellow. The M type are the coolest stars; they are characterized by temperatures of less than 3,500 K and a deep reddish colour. The O-M group is often supplemented by R, N, and S stars, which differ from the others in chemical composition.
Spectral analyses have provided much information about stellar composition. With several exceptions, such as ionized hydrogen in extremely hot stars, only the atoms of the more abundant constituent elements in a star are able to produce observable absorption lines. The spectrum of a G star, for instance, is dominated by the lines of such metals as iron, calcium, sodium, magnesium, and titanium.
The mass of a star is difficult to determine directly except in the case of visual binaries (i.e., stellar objects that can be seen as double stars in optical telescopes). Binaries, particularly those consisting of a dwarf and a supergiant star, provide the most extensive data on stellar dimensions. The angular diameters of supergiants were measured in the 1920s with the Michelson stellar interferometer. This instrument, based on the principle of light interference, provided an effective means of measuring the large angular diameters of bright stars, but it could not be used for stellar bodies of smaller apparent size. Several decades later, astronomers developed the photo-correlation interferometric method suitable for measuring the diameters of such stars. Since the late 1960s investigators have relied on another procedure, known as speckle interferometry, to reproduce the true disks of red supergiants.
Important generalizations concerning the nature and evolution of stars can be made from correlations between certain observed properties and from statistical results. The correlation between spectrum and absolute magnitude (or luminosity) is of particular significance. A graph called the Hertzsprung-Russell (H-R) diagram is generally used to show the absolute magnitudes of stars plotted against their spectral classes. The stars tend to cluster together in certain parts of the diagram. Most of them fall along a narrow band designated the main sequence that extends from the upper left of the diagram (hot, luminous stars) to its lower right (cool, dim stars). The stars of the main sequence, which includes the Sun, are the dwarf stars. Another relatively large number of stars congregate above the main sequence in the upper right portion of the diagram. These cool stars of high luminosity are the giants, which are about 100 times as bright as the Sun. Even more luminous stars, the supergiants, lie at the very top of the diagram. Below the main sequence at the lower left are the white dwarfs, stars with high surface temperatures but low luminosity.
A star forms when a dense interstellar cloud of hydrogen and dust grains collapses inward under the force of its own gravity. As the cloud condenses, its density and internal temperature increase until reaching incandescence with a faint red glow. At this stage the star shines by energy from its gravitational contraction. As its interior temperature rises to a few million degrees, deuterium (heavy hydrogen) is first destroyed, followed by the transformation of lithium, beryllium, and boron into helium. The temperature of the core region continues to increase until it reaches a level high enough to support thermonuclear reactions--the proton-proton chain or the carbon cycle. The star ceases to contract at this juncture, entering the main-sequence stage, where it remains for much of its life. As time passes, the star's chemical composition changes. The hydrogen in its core is converted into helium, and the central temperature gradually rises.
The change in composition is accompanied by changes in the star's structure, size, and luminosity. When all the hydrogen of the core has been exhausted and the central region consists almost entirely of inert helium, energy production begins to occur in a thin shell around the core. The core gradually increases in mass but shrinks in size because ever-increasing amounts of inert elements are fed in through the hydrogen-burning shell. The outer layers of the star expand significantly and cool, causing the star to become red in colour. At the same time, energy from the contracting core heats up the hydrogen in the surrounding region, which accelerates nuclear reactions and thereby increases luminosity.
The final stages of a star's evolution depend largely on two factors: (1) the mass of the stellar body, and (2) whether it is a component of a close binary system. A solitary star of less than 1.4 solar masses usually evolves from a red giant to a white dwarf when its outer layers drift away and leave behind a hot, compact core. This dense core constitutes the white dwarf star, and the expanding gaseous shell that temporarily surrounds it is known as a planetary nebula. When a lightweight star of a close binary system reaches the white-dwarf stage while its companion is becoming a red giant, it may acquire material from the latter's outer layers. This matter accumulates on the surface of the white dwarf and eventually triggers thermonuclear reactions. The star becomes a nova when the energy from these reactions blows off the accumulated matter in a brief but violent explosion.
Single stars that are more than five times as massive as the Sun can continue generating energy by fusion after having depleted their hydrogen supplies because their gravitational potential energy enables them to build up extremely high pressures and temperatures deep in their interior. In this way, they can successively create elements as heavy as iron in their cores. It is thought, however, that reactions involving iron fusion result in the catastrophic collapse of the stellar core. The outer layers of such a heavyweight star explode violently as a supernova, shining 109 times more brightly than a normal star for several months. Elements heavier than iron are created by neutron-capture reactions in the explosion itself, enriching the interstellar medium for future star-formation. After the supernova event, the core may remain as a neutron star. Composed of tightly packed neutrons, this type of star has a density many times greater than that of the Sun but a diameter of only about 20 km (12 miles). Many neutron stars emit radio radiation in the form of short pulses at very regular intervals. Such stellar objects are commonly called pulsars. If the remnant mass of the supernova is more than two or three solar masses, it is thought to continue collapsing inward and eventually to form a black hole, an object with so powerful a gravitational field that no form of matter or energy--not even light--can escape from it.
Excerpt from the Encyclopedia Britannica without permission.