A star is a large ball of hot gas, thousands to millions of kilometersin diameter, emitting large amounts of radiant energy from nuclearreactions in its interior. Stars differ fundamentally from planets in thatthey are self-luminous, whereas planets shine by reflected sunlight. Exceptfor the SUN, which is the nearest star, stars appear only as points oflight, even in the largest telescopes, because of their distance. The brightest stars have long been given names.
Most of the familiarnames originated with the ancient Greeks or with later Arab astronomers; anentirely different system was used by the Chinese, starting hundreds ofyears earlier, about 1000 BC. Polaris, the North Star, has a Greek name;Betelgeuse, a bright red star, has an Arabic name. Modern astronomersdesignate the bright stars according to the CONSTELLATIONS they are in. Thus, the brightest star in the Big Dipper (part of the constellation UrsaMajor) is called alpha Ursa Majoris.Order now
Polaris, in the Little Dipper (UrsaMinor), is gamma (designated by the Greek lower-case letter gamma) UrsaMinoris, and Betelgeuse, in Orion, is gamma Orionis. VARIABLE STARS (thosewhich periodically change in brightness) have lettered names, such as RRLyrae in the constellation Lyra. Fainter stars are known by their numbersin 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 intoclusters, and stars and star clusters are collected in the larger groupingscalled galaxies. Our GALAXY, the Milky Way, contains more than 100 billionstars.
Because tens of millions of other galaxies are known to exist, thetotal 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, yearafter year. They provided early astronomers with a reference system formeasuring the motions of planets (“wandering stars”), the Moon, and theSun. The westward rotation of the celestial sphere simply reflects thedaily eastward rotation of the Earth, and the Sun’s apparent motion amongthe stars reflects the Earth’s annual orbit around the Sun. As the construction of larger telescopes during the 19th centuryimproved the accuracy of determining stellar positions, it was found thatsome stars are not precisely “fixed. ” They move at various speeds, measuredas changes of direction in fractions of a second of arc per year, where onesecond of arc is the angular size of a pinhead 183 m (200 yd) away.
Most ofthe faint stars are truly fixed as viewed from Earth and are used as areference frame for the minute motions of nearby stars, known as PROPERMOTION. PARALLAX is another apparent motion of nearby stars. It is caused bythe 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) onone side of the Sun to 150 million km on the other side. Stellar parallaxcan be used to determine astronomical DISTANCE. If the shift is 1 second ofarc each way, the star is about 32 million million km (20 million millionmi) from an observer.
This distance is called the parsec and is equal to3. 26 light-years. The parallaxes of several thousand stars have beenmeasured during the past several decades. The nearest star is ProximaCentauri, at about 1 parsec (3.
3 light-years). Most of the measureddistances are greater than 20 parsecs (65 light-years), which shows why theaverage 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 inthe sky were described as “stars of the first magnitude. ” Later, themagnitude scale was defined more accurately: 6th magnitude stars are just1/100 as bright as 1st magnitude stars; 11th magnitude stars are 1/100 asbright 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 andblackness of images on a photographic plate exposed in a telescope-camera. With the photographic emulsions available in the early 1900s, a blue starthat appeared to the eye to have the same brightness as a red starphotographed much brighter. This discrepancy occurred because emulsions atthat time were much more sensitive to blue light than to red. Because ofthis variation, two magnitude scales came into use: visual magnitude andphotographic magnitude. The difference for any one star, photographicmagnitude minus visual magnitude, measures the color of that star–positivefor red stars, negative for blue (see COLOR INDEX).
By using filters andspecial emulsions, astronomers soon had several other magnitude scales,including ultraviolet and infrared. When photoelectric detectors wereintroduced, the brightnesses of stars were measured with a photoelectricphotometer at the focus of a telescope. Standard colors (wavelengths) oflight were adopted, and the symbols were changed to V and B, with U for theultraviolet scale, and several other letters for infrared scales. Measuring the brightness of a star on any of these scales iscomplicated by factors related to the Earth’s atmosphere, which absorbsmore light when a star is near the horizon than when it is overhead. Theatmosphere also absorbs different amounts of the different colors and canchange during the night because of changing dust or moisture in the air.
Nevertheless, by comparing a star with a standard at the same height abovethe 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 thetemperatures and energy output of stars, but it does not give the totalenergy output. Each measurement (U, B, V) gives only a fraction of thestar’s light reaching the Earth; even if the measurements are combined,they give only the part that is not absorbed as it passes through theEarth’s atmosphere. The atmosphere absorbs all light of short wavelengthsbelow ultraviolet and many of the long wavelengths above red.
A theoreticalcorrection 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 andspacecraft outside the Earth’s atmosphere. From parallax-distance measurements it is possible to calculate theabsolute bolometric magnitude, or luminosity, of a star, which is a measureof its brightness relative to the Sun if it were at the Sun’s distance froman 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 madeof an ethereal element different from matter on Earth. Their actualcomposition did not become known until the invention of the SPECTROSCOPE inthe 19th century.
Through the refraction of light by a prism (see PRISM,physics) or through its diffraction by a DIFFRACTION GRATING, the lightfrom a source is spread out into its different visual wavelengths, from redto blue; this is known as its SPECTRUM. The spectra of the Sun and starsexhibited bright and dark lines, which were shown to be caused by elementsemitting or absorbing light at specific wavelengths. Because each elementemits or absorbs light only at specific wavelengths, the chemicalcomposition of stars can be determined. In this way the spectroscopedemonstrated that the gases in the Sun and stars are those of commonelements such as hydrogen, helium, iron, and calcium at temperatures ofseveral thousand degrees. It was found that the average star’s atmosphereconsists mostly of hydrogen (87%) and helium (10%), an element discoveredfrom spectra of the Sun, with all other elements making up about 3%. Heliumactually was first discovered in the Sun’s spectrum.
At first, visual estimates of the strengths of spectral lines wereused to estimate the amounts of the elements present in the Sun and a fewstars, based on an analysis of the lines produced by a laboratory lightsource. When photographic emulsions came into use, the spectroscope becamethe spectrograph, with a photographic film or plate replacing the humaneye. During the first half of the 20th century, spectrographs were used ontelescopes to observe thousands of stars. On the spectrogram, theintensities of the lines are measured from the blackness of the film orplate. Most recently, photoelectric detectors are used to scan the spectrumin a spectrophotometer. Stellar spectra can also be measured byinterferometer techniques.
Although the ultraviolet, visual, and infrared parts of a star’sspectrum can be measured in this way, other techniques must be used, abovethe atmosphere, to measure the shorter wavelength spectra of X-ray starsand gamma-ray stars. Instead of gratings and prisms, various combinationsof filters and detectors are used to measure portions of the X-ray andgamma-ray spectra. At the other extreme (long wavelengths), radio spectraof stars and other radio sources are measured by “tuning” a radio telescopeto 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 amplifierat the focus. Radio spectra are much more accurate than optical spectra. Multiple radio telescopes, placed thousands of kilometers apart, candetermine the position of a radio-emitting star as accurately as an opticaltelescope 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 atHarvard University examined thousands of stellar spectra. Without concernfor the actual atmospheric gases or temperatures, Cannon classified eachspectrum 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 havebands, which are many lines close together, produced by molecules (seeHARVARD CLASSIFICATION OF STARS). Later studies showed that Cannon’sclasses 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’sformula, which gives the relative emissions of various colors from a hotbody. A cool star emits most of its light in the red; a hot star emits mostof its light in the blue. A measurement of the ratio of blue to red lightcoming from a star (its color index) determines its temperature. O starsare 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; andM stars have surface temperature =3,000 K. Other spectrographicmeasurements of absorption lines and emission lines help to confirm ormodify this so-called color temperature.
From 1911 to 1913, Einar Hertzsprung and H. N. Russell first plottedthe luminosity (L) versus the surface temperature (Ts) of stars, using as ameasure of temperature the spectral types determined by Cannon. TheHERTZSPRUNG-RUSSELL DIAGRAM first showed that highly luminous stars aremostly of classes O and B, with helium lines and surface temperature=25,000 K, whereas low-luminosity stars are mostly of class M and surfacetemperature =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 totalemission of radiant energy per unit area of a hot body’s surface at eachtemperature. From the bolometric luminosity, the total energy emitted isknown; from the temperature, the radiant energy emitted per squarecentimeter is known. The ratio gives the number of square centimeters, fromwhich the radius of the star can be calculated. This rough calculationshows that the radii of stars vary from 1/100 of that of the Sun for WHITEDWARFS to 400 times that of the Sun for SUPERGIANTS.
The radius of a nearbystar can also be measured directly with an interferometer on a telescope. Astronomers theorize that objects with a starlike composition but too smallto initiate nuclear reactions may also exist in the universe, helping toaccount for the “missing mass” of COSMOLOGY theories (see BROWN DWARF). Mass More than half of all stars are BINARY STARS–two or more stars thatorbit one another. About 100 orbits have been measured accurately.
Thesemeasurements provide perhaps the most important characteristic of a star:its mass. From Newton’s Laws of gravitation and motion, it is known thattwo highly massive stars must orbit (one around the other) faster than twostars of lesser mass at the same distance apart; thus the masses can becalculated from the orbit size and the period of the orbit. If the binarystars eclipse each other, this situation also gives estimates of eachstar’s diameter. Orbits of the planets show that the Sun’s mass is 2 X (10to the power of 33) g (2 billion billion billion tons, or about 333,000times 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 1/10the mass of the Sun. The mass of a star is also related to its luminosity; a high-mass starhas high luminosity, and a low-mass star has low luminosity.
TheMASS-LUMINOSITY RELATION states that the luminosity is approximatelyproportional to (mass) to the power of 3. 5. A star twice the mass of theSun 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, isclosely related to theories of stellar structure. In addition to luminosity and binary-star orbits, two systematicfeatures in the motions of stars relate to their masses. In many groups andclusters of stars, the stars have similar motions and similar Dopplershifts in the lines of their spectra (see RED SHIFT); these similaritiesare easy to pick out from the random motions of single stars. The smallermotions of stars within a cluster show the cluster’s total mass–the sum ofthe masses of all the stars bound together in it by their gravitation.
These internal motions can also be used statistically to determine thedistance from Earth to the cluster. More dramatic are the general motions of all the stars in the Sun’svicinity, showing a circulation around the center of the Milky Way Galaxy. Again, Newton’s laws apply, and from the average orbits of stars around thecenter, the mass of this GALAXY is found to be 100 billion times the Sun’smass. Because the orbital motions are faster near the center and slowerfarther away, individual motions can also be used to determine thedistances to individual stars. Since interstellar dust obscures more thanhalf of the stars in the Milky Way Galaxy, mass measurements give the onlyreliable 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 powerof 35)g.
Starspots Starspots (cooler regions on the surface of stars, similar to thefamiliar SUNSPOTS) are now known to exist on a number of relatively nearbystars. The disks of such stars can be mapped to some degree to show areasof differing temperature, using the technique known as speckleinterferometry (see INTERFEROMETER). The giant star Betelgeuse was observedin this manner as long ago as the mid-1970s. By means of spectral studies,astronomers have also been able to detect apparent granulation patterns onsome stars. Such patterns on the Sun are produced by convection, or therising and falling of hotter and cooler currents just below the visiblesurface. Analysis of stellar spectra to yield this kind of detail requiresthe use of supercomputers.
A larger, different kind of surface variation onstars 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 after1920, largely based on observations of the Sun. The photosphere is thevisible surface of a star and is the layer to which the surface temperatureand radius apply. Above the photosphere is an atmosphere, mostlytransparent, where gases absorb characteristic lines in the spectrum andreveal the chemical composition of the star. The temperature of the stellar atmosphere is lower than thetemperature of the photosphere. Above the atmosphere is a transparentCORONA of diffuse gas at high temperature.
For reasons as yet uncertain,outgoing energy from the Sun or star heats the corona to temperatures over1,000,000 K (1,800,000 deg F), so that it emits X rays of much shorterwavelength than visible light. The solar corona also has emission lines invisible light which give it the greenish glow visible during a total solareclipse. In the atmosphere and corona of a star, explosions known as flaresoccur in regions several thousand kilometers across, shooting outhigh-speed protons and electrons and causing plumes of higher temperaturein the corona. At a fairly constant rate, high-speed protons and electronsare also shot out in all directions to form the solar or stellar wind. TheSOLAR WIND has been detected by the two VOYAGER spacecraft and PIONEERS 10and 11 on their way out of the solar system.
Eventually they are expected tocross the outer boundary of the solar wind, the heliopause, whereinterstellar gas pressure stops the outflow of the wind. The knowledge of a star’s internal structure is almost entirelytheoretical, based on laboratory measurements of gases. Beneath thephotosphere are several layers, some where the hot, ionized gas isturbulent, and some where it is almost at rest. Calculations of structureare based on two principles: convective equilibrium, in which turbulencebrings the energy outward, and radiative equilibrium, in which radiationbrings the energy outward.
The temperature and density are calculated foreach 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 highenough (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 ofmass. Because the conversion of this mass (m) to energy (E) followsEinstein’s equation E = mcc (where c is the velocity of light), such areaction releases 6.
4 X (10 to the power of 18) ergs of energy per gram ofhydrogen, 60 million times more than chemical reactions such as the burningof hydrogen in oxygen. It is this enormous energy source that makeslong-lasting, self-luminous stars possible. In an attempt to determine the precise mechanism providing the energyfor stars, physicists in the early 1930s measured the rates of severalnuclear reactions in the laboratory. In 1938, Hans Bethe showed that thecarbon-nitrogen cycle could account for a star’s long-lasting luminosity(see CARBON CYCLE, astronomy). In Bethe’s theory, carbon acts as a catalystin the conversion of hydrogen to helium.
The small amount needed isconverted to nitrogen, then converted back to carbon to be used again. Thereaction rates at the temperature and density in the core of the Sun arefast enough to produce (10 to the power of 33) ergs/sec, the luminosity ofthe Sun. Later it was shown that the PROTON-PROTON REACTION could also producethe Sun’s luminosity. More recent studies show that in the Sun and smallerstars, where temperature and density in the core are lower than in largerstars, the proton-proton reaction beats out the Bethe cycle and can occurwith no carbon or nitrogen present, if the temperature is about 10,000,000K. In equations for the proton-proton reaction, the rates increase with thefourth power of the temperature, so that at a temperature of 20,000,000 Kthe rate is 16 times faster than at 10,000,000 K.
Lithium and beryllium areprobably also involved. The NEUTRINO is a very-low-mass particle that is produced in the Sun’score and can pass through its outer regions to enter space. One of thegreat mysteries of modern astrophysics is the failure of experiments todetect the neutrinos expected from nuclear reactions in the Sun. Whether by the Bethe cycle or by the proton-proton reaction, the Sunand other stars are converting hydrogen to helium in their cores at aconsiderable rate (600,000,000 tons/sec in the Sun).
Because helium hasdifferent characteristics, this conversion changes the structure of thestar. During the process there is a central core composed entirely ofhelium, a spherical shell around it in which hydrogen is being converted tohelium, and the rest of the star, composed mostly of hydrogen. When a largecore of helium has been created, the core may collapse, and new nuclearreactions may start as the temperature and density jump to very highvalues. When the temperature exceeds 100,000,000 K, helium is converted tocarbon by the triple-alpha (ionized helium) process.
Astrophysicists makeuse of the Hertzsprung-Russell diagram and large computers to calculate howstars evolve in this way. They find that stars of different masses evolvein different ways and at different rates. The most massive stars (ten timesthe Sun’s mass) rapidly change from blue giants to red giants and maybecome unstable and pulsate as variable stars during this stage. Stars oflesser mass, such as the Sun, spend a large fraction of their lives on themain sequence of the Hertzsprung-Russell diagram while they converthydrogen to helium. After several billion years, these stars become whitedwarfs. Depending on mass and other circumstances, a star may evolve to aNOVA or SUPERNOVA, PULSAR, NEUTRON STAR, or BLACK HOLE (see STELLAREVOLUTION).
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 FriendlyStars, rev. ed.
(1982); Goldberg, H. S. , and Scadron, M. D. , Physics ofStellar 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 STARSTABLE 1—————————————————————DistanceApparent BrightnessName(light-years)(magnitude)—————————————————————Sun – -26. 8Centauri A4. 3 -0. 01Centauri B4. 3 1.
33Centauri C4. 3 11. 05Barnard’s Star 5. 9 9.
54Wolf 359 7. 6 13. 53Lalande 21185 8. 1 7.
50Sirius A 8. 7 -1. 47Sirius B 8. 7 8. 68Luyten 726-8A 8. 9 12.
45Luyten 726-8B 8. 9 12. 95Ross 154 9. 4 10.
6Ross 24810. 3 12. 29Eridani 10. 7 3.
73Luyten 789-6 10. 8 12. 18Ross 12810. 8 11. 1061 Cygni A11. 2 5.
2261 Cygni B11. 2 6. 03Indi11. 2 4.
68Procyon A11. 3 0. 37Procyon B11. 3 10. 7—————————————————————SOURCE: Adapted from a table compiled by Alan H.
Batten in The Observer’sHandbook 1976 of the Royal Astronomical Society of Canada and a table Dramaof the Universe (1978) by George O. Abell (reprinted by permission of Holt,Rinehart and Winston). THE BRIGHTEST STARSTABLE 2————————————————————— Apparent BrightnessDistanceNameConstellation (magnitude)(light-year)—————————————————————Sun–26. 8 -Sirius ACanis Major -1. 47 8. 7Canopus Carina-0.
7298ArcturusBootes-0. 0636Centauri ACentaurus-0. 01 4. 3VegaLyra0. 0426.
5Capella Auriga 0. 0545RigelOrion 0. 14900Procyon ACanis Minor 0. 3711. 3BetelgeuseOrion 0. 41520AchernarEridanus0.
51118CentauriCentaurus0. 63490Altair Aquila 0. 7716. 5Crucis Crux0. 87400AldebaranTaurus 0. 8668SpicaVirgo 0.
91220Antares Scorpius0. 92520FomalhautPiscis Austrinus1. 1522. 6Pollux Gemini 1. 1635DenebCygnus 1. 261,600Crucis Crux1.
28490—————————————————————SOURCE: Adapted from a table compiled by Donald A. MacRae in The Observer’sHandbook 1976 of the Royal Astronomical Society of Canada and a table inContemporary Astronomy, 2d. , by Jay m. Pasachoff, Holt/Saunders, 1980.