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ФФФФФФФФФ>ФФФФФФФФФ>ФФФФФФФФФ>Chop Here>ФФФФФФФФФ>ФФФФФФФФФ>ФФФФФФФФФ>ФФФФФФФФФ A star is a large ball of hot gas, thousands to millions of kilometers in diameter, emitting large amounts of radiant energy from nuclear reactions in its interior. Stars differ fundamentally from planets in that they are self-luminous, whereas planets shine by reflected sunlight. Except for the SUN, which is the nearest star, stars appear only as points of light, even in the largest telescopes, because of their distance. The brightest stars have long been given names. Most of the familiar names originated with the ancient Greeks or with later Arab astronomers; an entirely different system was used by the Chinese, starting hundreds of years earlier, about 1000 BC. Polaris, the North Star, has a Greek name; Betelgeuse, a bright red star, has an Arabic name. Modern astronomers designate the bright stars according to the CONSTELLATIONS they are in. Thus, the brightest star in the Big Dipper (part of the constellation Ursa Major) is called alpha Ursa Majoris. Polaris, in the Little Dipper (Ursa Minor), is gamma (designated by the Greek lower-case letter gamma) Ursa Minoris, and Betelgeuse, in Orion, is gamma Orionis. VARIABLE STARS (those which periodically change in brightness) have lettered names, such as RR Lyrae in the constellation Lyra. Fainter stars are known by their numbers in a catalog; HD 12938 is the 12,938th star in the Henry Draper Catalogue. CHARACTERISTICS OF STARS Each star in the universe has its own position, motion, size, mass, chemical composition, and temperature. Some stars are grouped into clusters, and stars and star clusters are collected in the larger groupings called galaxies. Our GALAXY, the Milky Way, contains more than 100 billion stars. Because tens of millions of other galaxies are known to exist, the total number of stars in the universe exceeds a billion billion. Positions, Motions, and Distances Stars are seen in the same relative positions, night after night, year after year. They provided early astronomers with a reference system for measuring the motions of planets ("wandering stars"), the Moon, and the Sun. The westward rotation of the celestial sphere simply reflects the daily eastward rotation of the Earth, and the Sun's apparent motion among the stars reflects the Earth's annual orbit around the Sun. As the construction of larger telescopes during the 19th century improved the accuracy of determining stellar positions, it was found that some stars are not precisely "fixed." They move at various speeds, measured as changes of direction in fractions of a second of arc per year, where one second of arc is the angular size of a pinhead 183 m (200 yd) away. Most of the faint stars are truly fixed as viewed from Earth and are used as a reference frame for the minute motions of nearby stars, known as PROPER MOTION. PARALLAX is another apparent motion of nearby stars. It is caused by the Earth's orbit around the Sun: the star seems to shift, first one way, then the other, as the Earth moves from 150 million km (93 million mi) on one side of the Sun to 150 million km on the other side. Stellar parallax can be used to determine astronomical DISTANCE. If the shift is 1 second of arc each way, the star is about 32 million million km (20 million million mi) from an observer. This distance is called the parsec and is equal to 3.26 light-years. The parallaxes of several thousand stars have been measured during the past several decades. The nearest star is Proxima Centauri, at about 1 parsec (3.3 light-years). Most of the measured distances are greater than 20 parsecs (65 light-years), which shows why the average star in the sky is so much fainter than the nearby Sun. Brightness and Luminosity Star brightness was first estimated by eye, and the brightest stars in the sky were described as "stars of the first magnitude." Later, the magnitude scale was defined more accurately: 6th magnitude stars are just 1/100 as bright as 1st magnitude stars; 11th magnitude stars are 1/100 as bright as 6th magnitude, and so on. The magnitude scale is logarithmic; that is, each magnitude corresponds to a factor of 1/2.54, because (1/2.54) to the power of 5 =1/100 (see MAGNITUDE). Photographs are also used to measure star brightness from the size and blackness of images on a photographic plate exposed in a telescope-camera. With the photographic emulsions available in the early 1900s, a blue star that appeared to the eye to have the same brightness as a red star photographed much brighter. This discrepancy occurred because emulsions at that time were much more sensitive to blue light than to red. Because of this variation, two magnitude scales came into use: visual magnitude and photographic magnitude. The difference for any one star, photographic magnitude minus visual magnitude, measures the color of that star--positive for red stars, negative for blue (see COLOR INDEX). By using filters and special emulsions, astronomers soon had several other magnitude scales, including ultraviolet and infrared. When photoelectric detectors were introduced, the brightnesses of stars were measured with a photoelectric photometer at the focus of a telescope. Standard colors (wavelengths) of light were adopted, and the symbols were changed to V and B, with U for the ultraviolet scale, and several other letters for infrared scales. Measuring the brightness of a star on any of these scales is complicated by factors related to the Earth's atmosphere, which absorbs more light when a star is near the horizon than when it is overhead. The atmosphere also absorbs different amounts of the different colors and can change during the night because of changing dust or moisture in the air. Nevertheless, by comparing a star with a standard at the same height above the horizon, astronomers using photoelectric photometers can measure U, B, and V magnitudes with an accuracy of 0.01 magnitude (see PHOTOMETRY, ASTRONOMICAL). Such photometry has provided a great deal of information regarding the temperatures and energy output of stars, but it does not give the total energy output. Each measurement (U, B, V) gives only a fraction of the star's light reaching the Earth; even if the measurements are combined, they give only the part that is not absorbed as it passes through the Earth's atmosphere. The atmosphere absorbs all light of short wavelengths below ultraviolet and many of the long wavelengths above red. A theoretical correction can be made, based on the star's temperature, to give a "bolometric" magnitude, m(b), adding the energy absorbed by the atmosphere. True bolometric magnitudes, however, are measured only from rockets and spacecraft outside the Earth's atmosphere. From parallax-distance measurements it is possible to calculate the absolute bolometric magnitude, or luminosity, of a star, which is a measure of its brightness relative to the Sun if it were at the Sun's distance from an observer on Earth. During the 1920s it was found that some stars (giants) are 100,000 times as luminous as the Sun; others (white dwarfs) are 1,000 times less luminous. Composition During ancient times and the Middle Ages stars were thought to be made of an ethereal element different from matter on Earth. Their actual composition did not become known until the invention of the SPECTROSCOPE in the 19th century. Through the refraction of light by a prism (see PRISM, physics) or through its diffraction by a DIFFRACTION GRATING, the light from a source is spread out into its different visual wavelengths, from red to blue; this is known as its SPECTRUM. The spectra of the Sun and stars exhibited bright and dark lines, which were shown to be caused by elements emitting or absorbing light at specific wavelengths. Because each element emits or absorbs light only at specific wavelengths, the chemical composition of stars can be determined. In this way the spectroscope demonstrated that the gases in the Sun and stars are those of common elements such as hydrogen, helium, iron, and calcium at temperatures of several thousand degrees. It was found that the average star's atmosphere consists mostly of hydrogen (87%) and helium (10%), an element discovered from spectra of the Sun, with all other elements making up about 3%. Helium actually was first discovered in the Sun's spectrum. At first, visual estimates of the strengths of spectral lines were used to estimate the amounts of the elements present in the Sun and a few stars, based on an analysis of the lines produced by a laboratory light source. When photographic emulsions came into use, the spectroscope became the spectrograph, with a photographic film or plate replacing the human eye. During the first half of the 20th century, spectrographs were used on telescopes to observe thousands of stars. On the spectrogram, the intensities of the lines are measured from the blackness of the film or plate. Most recently, photoelectric detectors are used to scan the spectrum in a spectrophotometer. Stellar spectra can also be measured by interferometer techniques. Although the ultraviolet, visual, and infrared parts of a star's spectrum can be measured in this way, other techniques must be used, above the atmosphere, to measure the shorter wavelength spectra of X-ray stars and gamma-ray stars. Instead of gratings and prisms, various combinations of filters and detectors are used to measure portions of the X-ray and gamma-ray spectra. At the other extreme (long wavelengths), radio spectra of stars and other radio sources are measured by "tuning" a radio telescope to different frequencies. A radio telescope--the largest is more than 305 m (1,000 ft) across--is like a giant optical reflector with a radio amplifier at the focus. Radio spectra are much more accurate than optical spectra. Multiple radio telescopes, placed thousands of kilometers apart, can determine the position of a radio-emitting star as accurately as an optical telescope can, to better than 0.1 second of arc (see RADIO ASTRONOMY). Spectral Type and Surface Temperature During the early decades of the 20th century, Annie J. Cannon at Harvard University examined thousands of stellar spectra. Without concern for the actual atmospheric gases or temperatures, Cannon classified each spectrum as A, B, C, . . .S, depending on the number of absorption lines. Class A has few strong lines, class F has more, and classes M to S have bands, which are many lines close together, produced by molecules (see HARVARD CLASSIFICATION OF STARS). Later studies showed that Cannon's classes are a measure of surface temperature in the sequence O, B, A, F, G, K, M, R, N, S. This measurement is based partly on physicist Max Planck's formula, which gives the relative emissions of various colors from a hot body. A cool star emits most of its light in the red; a hot star emits most of its light in the blue. A measurement of the ratio of blue to red light coming from a star (its color index) determines its temperature. O stars are hot (surface temperature =30,000 K); A stars have surface temperature = 10,000 K; G stars, such as the Sun, have surface temperature =6,000 K; and M stars have surface temperature =3,000 K. Other spectrographic measurements of absorption lines and emission lines help to confirm or modify this so-called color temperature. From 1911 to 1913, Einar Hertzsprung and H. N. Russell first plotted the luminosity (L) versus the surface temperature (Ts) of stars, using as a measure of temperature the spectral types determined by Cannon. The HERTZSPRUNG-RUSSELL DIAGRAM first showed that highly luminous stars are mostly of classes O and B, with helium lines and surface temperature =25,000 K, whereas low-luminosity stars are mostly of class M and surface temperature =3,000 K. Size Once the temperature and the bolometric luminosity of a star are known, its size can easily be calculated. Planck's formula gives the total emission of radiant energy per unit area of a hot body's surface at each temperature. From the bolometric luminosity, the total energy emitted is known; from the temperature, the radiant energy emitted per square centimeter is known. The ratio gives the number of square centimeters, from which the radius of the star can be calculated. This rough calculation shows that the radii of stars vary from 1/100 of that of the Sun for WHITE DWARFS to 400 times that of the Sun for SUPERGIANTS. The radius of a nearby star can also be measured directly with an interferometer on a telescope. Astronomers theorize that objects with a starlike composition but too small to initiate nuclear reactions may also exist in the universe, helping to account for the "missing mass" of COSMOLOGY theories (see BROWN DWARF). Mass More than half of all stars are BINARY STARS--two or more stars that orbit one another. About 100 orbits have been measured accurately. These measurements provide perhaps the most important characteristic of a star: its mass. From Newton's Laws of gravitation and motion, it is known that two highly massive stars must orbit (one around the other) faster than two stars of lesser mass at the same distance apart; thus the masses can be calculated from the orbit size and the period of the orbit. If the binary stars eclipse each other, this situation also gives estimates of each star's diameter. Orbits of the planets show that the Sun's mass is 2 X (10 to the power of 33) g (2 billion billion billion tons, or about 333,000 times the Earth's mass). Orbits of binary stars show that some stars (giants) are 40 times the mass of the Sun, and others (dwarfs) only 1/10 the mass of the Sun. The mass of a star is also related to its luminosity; a high-mass star has high luminosity, and a low-mass star has low luminosity. The MASS-LUMINOSITY RELATION states that the luminosity is approximately proportional to (mass) to the power of 3.5. A star twice the mass of the Sun will have luminosity 2 to the power of 3.5, or 11.3 times the Sun's. This fact, together with the temperatures and compositions of stars, is closely related to theories of stellar structure. In addition to luminosity and binary-star orbits, two systematic features in the motions of stars relate to their masses. In many groups and clusters of stars, the stars have similar motions and similar Doppler shifts in the lines of their spectra (see RED SHIFT); these similarities are easy to pick out from the random motions of single stars. The smaller motions of stars within a cluster show the cluster's total mass--the sum of the masses of all the stars bound together in it by their gravitation. These internal motions can also be used statistically to determine the distance from Earth to the cluster. More dramatic are the general motions of all the stars in the Sun's vicinity, showing a circulation around the center of the Milky Way Galaxy. Again, Newton's laws apply, and from the average orbits of stars around the center, the mass of this GALAXY is found to be 100 billion times the Sun's mass. Because the orbital motions are faster near the center and slower farther away, individual motions can also be used to determine the distances to individual stars. Since interstellar dust obscures more than half of the stars in the Milky Way Galaxy, mass measurements give the only reliable estimate of the total number of stars in the Galaxy, 100 billion, each with a mass between (10 to the power of 32)g and 2 X (10 to the power of 35)g. Starspots Starspots (cooler regions on the surface of stars, similar to the familiar SUNSPOTS) are now known to exist on a number of relatively nearby stars. The disks of such stars can be mapped to some degree to show areas of differing temperature, using the technique known as speckle interferometry (see INTERFEROMETER). The giant star Betelgeuse was observed in this manner as long ago as the mid-1970s. By means of spectral studies, astronomers have also been able to detect apparent granulation patterns on some stars. Such patterns on the Sun are produced by convection, or the rising and falling of hotter and cooler currents just below the visible surface. Analysis of stellar spectra to yield this kind of detail requires the use of supercomputers. A larger, different kind of surface variation on stars has been reported by some astronomers, who call these variations "starpatches." STRUCTURE OF STARS The structure of a typical star was worked out by astrophysicists after 1920, largely based on observations of the Sun. The photosphere is the visible surface of a star and is the layer to which the surface temperature and radius apply. Above the photosphere is an atmosphere, mostly transparent, where gases absorb characteristic lines in the spectrum and reveal the chemical composition of the star. The temperature of the stellar atmosphere is lower than the temperature of the photosphere. Above the atmosphere is a transparent CORONA of diffuse gas at high temperature. For reasons as yet uncertain, outgoing energy from the Sun or star heats the corona to temperatures over 1,000,000 K (1,800,000 deg F), so that it emits X rays of much shorter wavelength than visible light. The solar corona also has emission lines in visible light which give it the greenish glow visible during a total solar eclipse. In the atmosphere and corona of a star, explosions known as flares occur in regions several thousand kilometers across, shooting out high-speed protons and electrons and causing plumes of higher temperature in the corona. At a fairly constant rate, high-speed protons and electrons are also shot out in all directions to form the solar or stellar wind. The SOLAR WIND has been detected by the two VOYAGER spacecraft and PIONEERS 10 and 11 on their way out of the solar system.Eventually they are expected to cross the outer boundary of the solar wind, the heliopause, where interstellar gas pressure stops the outflow of the wind. The knowledge of a star's internal structure is almost entirely theoretical, based on laboratory measurements of gases. Beneath the photosphere are several layers, some where the hot, ionized gas is turbulent, and some where it is almost at rest. Calculations of structure are based on two principles: convective equilibrium, in which turbulence brings the energy outward, and radiative equilibrium, in which radiation brings the energy outward. The temperature and density are calculated for each depth, using the characteristics of the mix of gases (hydrogen, helium, and heavier elements) derived from the spectrum of the atmosphere. The pressure is calculated from the weight of the gases overhead. Eventually, deep in the interior the temperature and density are high enough (10,000,000 K and 30 g/cu cm) for a nuclear reaction to occur, converting four hydrogen atoms to one helium atom, with a 0.7% loss of mass. Because the conversion of this mass (m) to energy (E) follows Einstein's equation E = mcc (where c is the velocity of light), such a reaction releases 6.4 X (10 to the power of 18) ergs of energy per gram of hydrogen, 60 million times more than chemical reactions such as the burning of hydrogen in oxygen. It is this enormous energy source that makes long-lasting, self-luminous stars possible. In an attempt to determine the precise mechanism providing the energy for stars, physicists in the early 1930s measured the rates of several nuclear reactions in the laboratory. In 1938, Hans Bethe showed that the carbon-nitrogen cycle could account for a star's long-lasting luminosity (see CARBON CYCLE, astronomy). In Bethe's theory, carbon acts as a catalyst in the conversion of hydrogen to helium. The small amount needed is converted to nitrogen, then converted back to carbon to be used again. The reaction rates at the temperature and density in the core of the Sun are fast enough to produce (10 to the power of 33) ergs/sec, the luminosity of the Sun. Later it was shown that the PROTON-PROTON REACTION could also produce the Sun's luminosity. More recent studies show that in the Sun and smaller stars, where temperature and density in the core are lower than in larger stars, the proton-proton reaction beats out the Bethe cycle and can occur with no carbon or nitrogen present, if the temperature is about 10,000,000 K. In equations for the proton-proton reaction, the rates increase with the fourth power of the temperature, so that at a temperature of 20,000,000 K the rate is 16 times faster than at 10,000,000 K. Lithium and beryllium are probably also involved. The NEUTRINO is a very-low-mass particle that is produced in the Sun's core and can pass through its outer regions to enter space. One of the great mysteries of modern astrophysics is the failure of experiments to detect the neutrinos expected from nuclear reactions in the Sun. Whether by the Bethe cycle or by the proton-proton reaction, the Sun and other stars are converting hydrogen to helium in their cores at a considerable rate (600,000,000 tons/sec in the Sun). Because helium has different characteristics, this conversion changes the structure of the star. During the process there is a central core composed entirely of helium, a spherical shell around it in which hydrogen is being converted to helium, and the rest of the star, composed mostly of hydrogen. When a large core of helium has been created, the core may collapse, and new nuclear reactions may start as the temperature and density jump to very high values. When the temperature exceeds 100,000,000 K, helium is converted to carbon by the triple-alpha (ionized helium) process. Astrophysicists make use of the Hertzsprung-Russell diagram and large computers to calculate how stars evolve in this way. They find that stars of different masses evolve in different ways and at different rates. The most massive stars (ten times the Sun's mass) rapidly change from blue giants to red giants and may become unstable and pulsate as variable stars during this stage. Stars of lesser mass, such as the Sun, spend a large fraction of their lives on the main sequence of the Hertzsprung-Russell diagram while they convert hydrogen to helium. After several billion years, these stars become white dwarfs. Depending on mass and other circumstances, a star may evolve to a NOVA or SUPERNOVA, PULSAR, NEUTRON STAR, or BLACK HOLE (see STELLAR EVOLUTION). Bibliography: Barrow, J. D., and Silk, Joseph, The Left Hand of Creation (1983); Abell, G., Exploration of the Universe (1969); Baade, Walter, Evolution of Stars and Galaxies (1975); Evans Martin, Martha, The Friendly Stars, rev. ed. (1982); Goldberg, H. S., and Scadron, M. D., Physics of Stellar Evolution and Cosmology (1982); Hall, Douglas, "Starspots," Astronomy, February 1983; Kruse, W., and Dieckvoss, W., The Stars (1957); Kyselka, Will, and Lanterman, Ray, North Star to Southern Cross (1976); Meadows, A. J., Stellar Evolution (1978); Page, Thornton, and Page, L. W., Starlight (1967) and Stars and Clouds of the Milky Way (1968); Shklovskii, Iosif S., Stars: Their Birth, Life and Death, trans. by Richard Rodman (1978). THE NEAREST STARS TABLE 1 --------------------------------------------------------------- Distance Apparent Brightness Name (light-years) (magnitude) --------------------------------------------------------------- Sun - -26.8 Centauri A 4.3 -0.01 Centauri B 4.3 1.33 Centauri C 4.3 11.05 Barnard's Star 5.9 9.54 Wolf 359 7.6 13.53 Lalande 21185 8.1 7.50 Sirius A 8.7 -1.47 Sirius B 8.7 8.68 Luyten 726-8A 8.9 12.45 Luyten 726-8B 8.9 12.95 Ross 154 9.4 10.6 Ross 248 10.3 12.29 Eridani 10.7 3.73 Luyten 789-6 10.8 12.18 Ross 128 10.8 11.10 61 Cygni A 11.2 5.22 61 Cygni B 11.2 6.03 Indi 11.2 4.68 Procyon A 11.3 0.37 Procyon B 11.3 10.7 --------------------------------------------------------------- SOURCE: Adapted from a table compiled by Alan H. Batten in The Observer's Handbook 1976 of the Royal Astronomical Society of Canada and a table Drama of the Universe (1978) by George O. Abell (reprinted by permission of Holt, Rinehart and Winston). THE BRIGHTEST STARS TABLE 2 --------------------------------------------------------------- Apparent Brightness Distance Name Constellation (magnitude) (light-year) --------------------------------------------------------------- Sun - -26.8 - Sirius A Canis Major -1.47 8.7 Canopus Carina -0.72 98 Arcturus Bootes -0.06 36 Centauri A Centaurus -0.01 4.3 Vega Lyra 0.04 26.5 Capella Auriga 0.05 45 Rigel Orion 0.14 900 Procyon A Canis Minor 0.37 11.3 Betelgeuse Orion 0.41 520 Achernar Eridanus 0.51 118 Centauri Centaurus 0.63 490 Altair Aquila 0.77 16.5 Crucis Crux 0.87 400 Aldebaran Taurus 0.86 68 Spica Virgo 0.91 220 Antares Scorpius 0.92 520 Fomalhaut Piscis Austrinus 1.15 22.6 Pollux Gemini 1.16 35 Deneb Cygnus 1.26 1,600 Crucis Crux 1.28 490 --------------------------------------------------------------- SOURCE: Adapted from a table compiled by Donald A. MacRae in The Observer's Handbook 1976 of the Royal Astronomical Society of Canada and a table in Contemporary Astronomy, 2d., by Jay m. Pasachoff, Holt/Saunders, 1980.