Sun, river, United States
Sun, river, c.130 mi (210 km) long, rising in the Rocky Mts., NW Mont., and flowing generally E to the Missouri River at Great Falls. The Sun River project of the U.S. Bureau of Reclamation utilizes the Sun and its tributaries to irrigate c.92,000 acres (37,230 hectares) of land. Of the system of dams and reservoirs, Gibson Dam is one of the project's largest.sun, in astronomy
sun, intensely hot, self-luminous body of gases at the center of the solar system solar system, the sun and the surrounding planets, natural satellites, dwarf planets, asteroids, meteoroids, and comets that are bound by its gravity. The sun is by far the most massive part of the solar system, containing almost 99.9% of the system's total mass...... Click the link for more information. . Its gravitational attraction maintains the planets, comets, and other bodies of the solar system in their orbits.
General Characteristics of the Sun
The sun is actually a star of about medium size; it appears larger than the other stars because of its relative nearness to the earth. The earth's distance from the sun varies from 91,377,000 mi (147,053,000 km) at perihelion to 94,537,000 mi (152,138,000 km) at aphelion (see apsis apsis (pl. apsides), point in the orbit of a body where the body is neither approaching nor receding from another body about which it revolves. Any elliptical orbit has two apsides.
..... Click the link for more information. ). The mean distance is c.92,960,000 mi (149,591,000 km); this is taken as the astronomical unit astronomical unit (AU), mean distance between the earth and sun; one AU is c.92,960,000 mi (149,604,970 km). The astronomical unit is the principal unit of measurement within the solar system, e.g., Mercury is just over 1-3 AU and Pluto is about 39 AU from the sun.
..... Click the link for more information. (AU) of distance used for measuring distances within the solar system. The sun is approximately 865,400 mi (1,392,000 km) in diameter, and its volume is about 1,300,000 times that of the earth. Its mass is almost 700 times the total mass of all the bodies in the solar system and 332,000 times that of the earth. The sun's surface gravity is almost 28 times that of the earth; i.e., a body on the surface of the sun would weigh about 28 times its weight on earth. The density of the material composing the sun is about one fourth that of the earth; compared with water, the sun's average density is 1.41. At its center, the sun has a density of over 100 times that of water, a temperature of 10 to 20 million degrees Celsius, and a pressure of over 1 billion atmospheres.
Observations of sunspots and studies of the solar spectrum indicate that the sun rotates on its axis from east to west; because of its gaseous nature its rate of rotation varies somewhat with latitude, the speed being greatest (a period of almost 25 days) in the equatorial region and least at the poles (a period of about 35 days). The axis of the sun is inclined at an angle of about 7° to the plane of the ecliptic ecliptic , the great circle on the celestial sphere that lies in the plane of the earth's orbit (called the plane of the ecliptic). Because of the earth's yearly revolution around the sun, the sun appears to move in an annual journey through the heavens with the
..... Click the link for more information. .
The bright surface of the sun is called the photosphere photosphere, luminous, apparently opaque layer of gases that forms the visible surface of the sun or any other star. The photosphere lies between the dense interior gases and the more attenuated gases of the chromosphere.
..... Click the link for more information. . Its temperature is about 6,000°C;. The photosphere appears darker near the edge (limb) of the sun's disk because of greater absorption of light by the sun's atmosphere in this area; this phenomenon is called limb darkening. During an eclipse of the sun the chromosphere chromosphere [Gr.,=color sphere], layer of rarefied, transparent gases in the solar atmosphere; it measures 6,000 mi (9,700 km) in thickness and lies between the photosphere (the sun's visible surface) and the corona (its outer atmosphere).
..... Click the link for more information. and the corona corona, luminous envelope surrounding the sun, outside the chromosphere. Its density is less than one billionth that of the earth's atmosphere. The corona is visible only at the time of totality during a total eclipse of the sun.
..... Click the link for more information. (the outer layers of the sun's atmosphere) are observed. Also of interest is the high-speed, tenuous extension of the corona known as the solar wind solar wind, stream of ionized hydrogen—protons and electrons—with an 8% component of helium ions and trace amounts of heavier ions that radiates outward from the sun at high speeds.
..... Click the link for more information. .
Production of Solar Energy
The vast and continual production of solar energy cannot be attributed merely to combustion, to the gradual cooling of a hot body, to the fall of meteorites into the sun, or to gradual shrinkage with transformation of potential energy into heat (a theory proposed by Helmholtz). The theory of relativity with its implication of the equivalence of mass and energy led to the assumption that energy stored in the atoms constituting the sun's gases is constantly being released by conversion of some of the masses of the atom's nuclei during nuclear transmutations (see nuclear energy nuclear energy, the energy stored in the nucleus of an atom and released through fission, fusion, or radioactivity. In these processes a small amount of mass is converted to energy according to the relationship E = mc2, where E
..... Click the link for more information. ). H. A. Bethe proposed a cycle of nuclear reactions known as the carbon cycle, or CNO bi-cycle, to account for the nuclear changes. In this cycle carbon acts much as a catalyst, while hydrogen is transformed by a series of reactions into helium and large amounts of high-energy gamma radiation are released. It is now thought that the so-called proton-proton process is a more important energy source; this process begins with the collision of two protons and ends with the production of helium, while gamma radiation is released throughout.
See nucleosynthesis nucleosynthesis or nucleogenesis, in astronomy, production of all the chemical elements from the simplest element, hydrogen, by thermonuclear reactions within stars, supernovas, and in the big bang at the beginning of the universe (see nucleus;
..... Click the link for more information. ; stellar evolution stellar evolution, life history of a star, beginning with its condensation out of the interstellar gas (see interstellar matter) and ending, sometimes catastrophically, when the star has exhausted its nuclear fuel or can no longer adjust itself to a stable
..... Click the link for more information. .
The Study of the Sun
By means of the spectroscope spectroscope, optical instrument for producing spectral lines and measuring their wavelengths and intensities, used in spectral analysis (see spectrum). When a material is heated to incandescence it emits light that is characteristic of the atomic makeup of the
..... Click the link for more information. much has been learned about the composition of the sun. There are numerous dark lines of varying widths in the solar spectrum spectrum, arrangement or display of light or other form of radiation separated according to wavelength, frequency, energy, or some other property. Beams of charged particles can be separated into a spectrum according to mass in a mass spectrometer (see mass
..... Click the link for more information. . These were first intensively studied by Joseph Fraunhofer and are commonly known by his name. From a study of the lines the chemical composition of the sun is determined on the basis of the discovery by Kirchhoff that the dark lines correspond in position to the bright lines characteristic of the spectra produced by elements in the laboratory. The darkness of the lines in the sun's spectrum is attributed to the presence of a slightly cooler layer of gases above the photosphere, known as the reversing layer, which absorbs selectively the light of the photosphere and thus causes dark lines instead of bright ones to be observed through the spectroscope. By comparison of the sun's spectrum with laboratory spectra of incandescent elements, most of the elements known on earth have been identified in the sun's atmosphere.
Beyond the red portion of the visible solar spectrum is the infrared spectrum; for the study of these heat rays S. P. Langley invented the bolometer, a highly sensitive electrical device for measuring temperature. Solar heat and energy are measured by an instrument called the pyrheliometer. Other instruments devised especially for the study of the sun are the coronagraph and the spectroheliograph spectroheliograph, device for photographing the surface of the sun in a single wavelength of light, usually one corresponding to a chief element contained in the sun, e.g., hydrogen or calcium; the resulting photograph is called a spectroheliogram.
..... Click the link for more information. . These instruments have revealed a number of interesting phenomena occurring during the periods of solar activity associated with sunspots sunspots, dark, usually irregularly shaped spots on the sun's surface that are actually solar magnetic storms. The Chinese recorded dark features on the sun seen with the naked eye in 28 B.C.
..... Click the link for more information. , e.g., faculae, plages (flocculi), prominences, and flares.
Importance to Terrestrial Life
Without the heat and light of the sun, life as we know it could not exist on the earth. Since solar energy is used by green plants in the process of photosynthesis, the sun is the ultimate source of the energy stored both in food and fossil fuels. Solar heating sets up convection currents, and thus is the source of the energy of moving air. Falling rain also owes its energy to the sun because of the relation of solar radiation to the water cycle.
Bibliography
See K. Hufbauer, Exploring the Sun: Solar Science since Galileo (1993); R. Krippenhahn, Discovering the Secrets of the Sun (1994); K. J. H. Phillips, Guide to the Sun (1995); P. O. Taylor, Beginners Guide to the Sun (1996); S. T. Suess and B. T. Tsurutani, ed., From the Sun: Auroras, Magnetic Storms, Solar Flares, Cosmic Rays (1998).
Sun
Star around which the components of the solar system revolve. It is about five billion years old and is the dominant body of the system, with more than 99% of its mass. It converts five million tons of matter into energy every second by nuclear fusion reactions in its core, producing neutrinos (see solar neutrino problem) and solar radiation. The small amount of this energy that penetrates Earth’s atmosphere provides the light and heat that support life. A sphere of luminous gas 864,950 mi (1,392,000 km) in diameter, the Sun has about 330,000 times the mass of Earth. Its core temperature is close to 27 million °F (15 million °C) and its surface temperature about 10,000 °F (6,000 °C). The Sun, a spectral type G (yellow) star, has fairly average properties for a main-sequence star (see Hertzsprung-Russell diagram). It rotates at different rates at different latitudes; one rotation takes 36 days at the poles but only 25 days at the equator. The visible surface, or photosphere, is in constant motion, with the number and position of sunspots changing in a regular solar cycle. External phenomena include magnetic activity extending into the chromosphere and corona, solar flares, solar prominences, and the solar wind. Effects on Earth include auroras and disruption of radio communications and power-transmission lines. Despite its activity, the Sun appears to have remained relatively unchanged for billions of years. See also eclipse; heliopause.
It all began in 1981 when Bavarian-born Andreas Bechtolsheim was licensing rights to a computer he designed. Named Sun for Stanford University Network and using off-the-shelf parts, it was an affordable workstation for engineers and scientists. In that year, he met Vinod Khosla, a native of India, who convinced him to form a company and expand. Khosla, Bechtolsheim and Scott McNealy, all Stanford MBAs, founded Sun in 1982.
Its first computers, the Sun-1 and subsequent Sun-2 were instant successes in the university market. Sun began to compete against its rival Apollo Computer, an east-coast workstation company, eventually surpassing it in sales (Apollo was later purchased by HP).
Sun has been a major force in open systems. Its computers have always run under Unix, which was licensed from AT&T and then later purchased outright. Sun and AT&T had formed such a tight alliance for a while that a host of Unix vendors formed the Open Software Foundation (OSF) in 1988 to keep Sun from dominating Unix.
In 1984, Bill Joy, head of R&D, designed NFS, which was broadly licensed and became the industry standard for file sharing. Sun later packaged its Unix components into a complete environment named Solaris, which it later ported to other platforms, including the Intel x86.
Sun used the Motorola 68K CPUs in its products until it designed its own RISC-based SPARC chips, which it launched with the SPARCstation 1 in 1989. Having gone through many iterations, SPARC CPUs are also made by Fujitsu and other third parties via licensing arrangements (see SPARC).
In the mid-1990s, Sun introduced the Java programming language and ushered in a new era for application development on the Internet (see Java and J2EE). See network computer and Sun-Netscape Alliance.
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| The Founders |
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| From left to right: Vinod Khosla, Bill Joy, Andreas Bechtolsheim and Scott McNealy. Although Joy was not a founder, he was hired shortly thereafter and became one of Sun's major contributors. (Image courtesy of Sun Microsystems, Inc.) |
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| The Sun-1 |
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| Sun's workstations were an instant success primarily in the university market. This led many professionals to the company who helped it grow steadily. (Image courtesy of Sun Microsystems, Inc.) |
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| The First SPARCstation |
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| In 1989, Sun introduced the SPARCstation 1, the first Sun computer that used the SPARC chip. (Image courtesy of Sun Microsystems, Inc.) |
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| Sun JavaStation |
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| Sun developed the Java programming language, and its network computer is aptly named. The JavaStation conforms to the NC Reference Profile and is available in several models. (Image courtesy of Sun Microsystems, Inc.) |
sun
www.solarviews.com/eng/sun.htm
www.michielb.nl/sun/kaft.htm
www.hao.ucar.edu/public/education/education.html
sun [sən]
Sun
| Sun - Sun Microsystems |
Sun
a hot plasma sphere that is the central body of the solar system. The sun is the closest star to the earth.
The mass of the sun is 1.990 × 1030 kg, which is 332,958 times greater than the mass of the earth. The sun contains 99.866 percent of the mass of the solar system. The solar parallax—that is, the angle subtended at the sun’s center by the earth’s equatorial radius when the earth is at its mean distance from the sun—is equal to 8.794“, or 4.263 × 10–5 radian (rad). The distance of the earth from the sun varies from 1.4710 × 1011 m in January to 1.5210 × 1011 m in July; the mean distance, which is known as the astronomical unit, is 1.4960 × 1011 m. The mean angular diameter of the sun is 1,919.26”, or 9.305 × 10–3 rad. The diameter of the sun is 1.392 × 109 m and is thus 109 times greater than the equatorial diameter of the earth. The mean density of the sun is 1.41 × 103 kg/m3. The gravitational acceleration at the sun’s surface is 273.98 m/sec2, and the escape velocity 6.18 × 105 m/sec. The effective surface temperature of the sun, as calculated in accordance with the Stefan-Boltzmann law from the total energy radiated by the sun, is 5770°K (see).
History of telescopic observation. The first telescopic observations of the sun were made by Galileo in 1611. He discovered sunspots and showed that the sun rotates about its axis. In 1843 the German astronomer H. Schwabe discovered the cyclicity of solar activity. The development of methods of spectrum analysis permitted study of the physical conditions on the sun. J. von Fraunhofer’s discovery of dark absorption lines in the solar spectrum in 1814 laid the foundation for study of the chemical composition of the sun. Observations of solar eclipses have been conducted regularly since 1836; results of these observations include the discovery of the solar corona, the solar chromosphere, and solar prominences. In 1913 the American astronomer G. Hale detected the Zeeman splitting of some Fraunhofer lines in sun-spot spectra and thereby demonstrated the existence of magnetic fields on the sun. By 1942, the Swedish astronomer B. Edlén and others had identified several lines of the coronal spectrum as lines of highly ionized elements; evidence was thus provided for a high temperature in the corona. In 1931, B. Lyot invented the corona-graph, which made possible observation of the corona and chromosphere at times other than during solar eclipses. The radio emission of the sun was discovered in the early 1940’s.
The development of magnetohydrodynamics and plasma physics in the second half of the 20th century gave a substantial impetus to the development of solar physics. The onset of the space age brought the application of the methods of space astronomy, or extraterrestrial astronomy, to the study of solar radiation in the ultraviolet and X-ray regions of the spectrum. These methods involve the use of rockets, automatic orbiting observatories on artificial earth satellites, and manned space laboratories.
In the USSR, solar research is carried on at the Crimean and Pulkovo observatories and at astronomical institutions in Moscow, Kiev, Tashkent, Alma-Ata, Abastumani, and Irkutsk. Most astrophysical observatories outside the USSR engage in solar research (seeASTRONOMICAL OBSERVATORIES AND INSTITUTES).
Rotation. The sun rotates about its axis in the same direction as the earth. The solar equator is inclined at 7°15’ to the plane of the ecliptic, that is, the plane of the earth’s orbit. The rotational speed can be determined from the apparent motion of various features of the solar atmosphere and from the shift of the lines of the spectrum of the limb, or edge, of the solar disk owing to the Doppler effect. Through these methods, the sun’s period of rotation was found to be different at different latitudes.
The positions of features on the solar surface are determined in terms of heliographic coordinates. The heliographic latitude is measured from the solar equator. The heliographic longitude is measured from the central meridian of the visible solar disk or from some other meridian selected as the prime meridian. It is assumed here that the sun rotates as a solid body. Astronomical yearbooks give the position of the prime meridian for each day. The yearbooks also give the position of the solar axis on the celestial sphere.
A point whose heliographic latitude is 17° completes a single rotation, relative to the earth, in 27.275 days; the period of rotation measured in this way is known as the synodic period. The sidereal period of rotation, that is, the period relative to the stars, at the same latitude is 25.38 days. The sidereal angular speed of rotation ω varies with the latitude φ in accordance with the equation (ω = 14.44° – 3°(sin2 φ) per day. The linear speed of rotation at the solar equator is approximately 2,000 m/sec.
The sun as a star. The sun is a typical yellow dwarf and lies near the middle of the main sequence of the Hertzsprung-Russell diagram. The apparent photovisual magnitude of the sun is –26.74, and the absolute visual magnitude Mv is +4.83. The color index of the sun for the blue (B) and visual (V) spectral regions is MB – Mv = 0.65. The sun’s spectral type is G2V. The speed of the sun in space relative to the nearest stars is 19.7 × 103 m/sec. The sun is located within one of the spiral arms of the Milky Way Galaxy at a distance of approximately 10 kiloparsecs from the Galaxy’s center. The sun’s period of revolution about the center of the Galaxy is approximately 200 million years. The age of the sun is approximately 5 × 109 years.
Internal structure. The internal structure of the sun is determined on the assumption that the sun is a spherically symmetric body and is in equilibrium. A model of the sun’s internal structure can be constructed on the basis of the following: the energy transport equation; the law of conservation of energy; the equation of state of an ideal gas; the Stefan-Boltzmann law; the conditions for hydrostatic, radiative, and convective equilibrium; ob-servationally determined values for the total radiation, total mass, and radius of the sun; and data on the sun’s chemical composition. Hydrogen is presumed to make up about 70 percent of the sun’s mass, and helium about 27 percent; the content of all other elements is presumed to be about 2.5 percent. On the basis of these assumptions, it has been calculated that the temperature at the sun’s center is (10–15) × 106°K, the density is about 1.5 × 105 kg/m3, and the pressure is 3.4 × 1016 newtons per m2(N/m2), or about 3 × 1016 atmospheres.
Nuclear reactions in the sun’s interior are presumed to be the energy source that replenishes the energy lost to radiation and maintains the high temperature of the sun. The mean quantity of energy generated within the sun is 1.92 ergs/g/sec. The release of energy is due to nuclear reactions in which hydrogen is converted into helium. Two series of such thermonuclear reactions are possible in the sun: the proton-proton chain, or hydrogen cycle, and the carbon cycle. It is most likely that the proton-proton chain predominates in the sun. This cycle consists of three reactions: (1) nuclei of deuterium, which is a heavy isotope of hydrogen of mass number 2, form from hydrogen nuclei; (2) nuclei of the helium isotope of mass number 3 form from the deuterium nuclei; and (3) nuclei of the stable isotope of helium of mass number 4 form from the light helium nuclei.
Energy transport. Energy is transported outward from the inner layers of the sun mainly through the absorption of electromagnetic radiation from lower layers and the subsequent reemission of the radiation. Since the temperature decreases with increasing distance from the center of the sun, the wavelength of the radiation transporting the greater part of the energy to the upper layers gradually increases (seeWIEN’S LAW OF RADIATION).
Convective transport of energy—that is, transport through the motion of hot material upward from inner layers and the motion of cooler material toward the center—plays an important role in the relatively high layers forming the convective zone of the sun. This zone begins at a depth of the order of 0.2 solar radius and is about 108 m thick. The speed of the convective motions increases with increasing distance from the center of the sun and in the outer portion of the convective zone reaches 2,000–2,500 m/sec.
In even higher layers—that is, in the solar atmosphere—energy is again transported through radiation. Some energy is carried to the upper layers of the solar atmosphere (the chromosphere and corona) by mechanical and magnetohydrodynamic waves that are generated in the convective zone but are absorbed only in these layers. The density of the upper atmosphere is very low, and the required removal of energy by radiation and thermal conduction is possible only if the kinetic temperature of these layers is sufficiently high. Finally, in the upper portion of the corona, most of the energy is carried away by the solar wind, which is an outflow of matter from the sun.
The temperature in each layer is at such a level that an energy balance is automatically maintained: the amount of energy acquired owing to absorption of radiation, to thermal conduction, or to convection is equal to the sum of the energy losses of the layer.
Total radiation. The total solar radiation can be determined from the irradiance from the sun at the earth’s surface. The average illuminance at the earth’s surface is about 100,000 lux when the sun is at the zenith. Outside the atmosphere, at the earth’s mean distance from the sun, the illuminance is 127,000 lux. The luminous intensity of the sun is 2.84 × 1027 international candles. The important quantity known as the solar constant is the total radiation energy received from the sun per min per cm2 on a surface that is outside the atmosphere at the earth’s mean distance from the sun and is perpendicular to the sun’s rays. The total energy output of the sun is 3.83 × 1026 watts (W), of which about 2 × 1017 W reach the earth. The mean luminance of the solar surface when observed outside the earth’s atmosphere is 1.98 × 109nits (nt), and the luminance of the center of the solar disk is 2.48 × 109 nt. The luminance of the solar disk decreases from the center of the disk to the limb. This decrease depends on the wavelength—for example, the luminance at the limb is about 0.2 of that of the center for light of wavelength 3,600 angstroms (Å) and is about 0.3 of the luminance of the center for light of wavelength 5,000 A. At the very edge of the solar disk, the luminance decreases by a factor of 100 over less than 1 second of arc. For this reason, the boundary of the solar disk appears very sharp.
Spectral composition of emitted light. The spectral composition of the light emitted by the sun—that is, the spectral energy distribution—roughly corresponds to the energy distribution in the radiation of a blackbody at a temperature of about 6000°K (after allowance is made for absorption in the earth’s atmosphere and for the Fraunhofer lines). Marked deviations, however, are observed in some parts of the spectrum. The energy maximum in the solar spectrum corresponds to the wavelength 4,600 Å.
The solar spectrum is a continuum on which more than 20,000 absorption lines (Fraunhofer lines) are superposed. More than 60 percent of the Fraunhofer lines have been identified as the spectral lines of known chemical elements by comparison of the wavelengths and relative intensities of the Fraunhofer lines with laboratory spectra. Study of the Fraunhofer lines gives information not only on the chemical composition of the solar atmosphere but also on the physical conditions in the layers in which particular lines arise.
The predominant element in the sun is hydrogen. The number of hydrogen atoms is four or five times greater than the number of helium atoms and is at least 1,000 times greater than the number of atoms of all the other elements combined. Among these other elements, the most abundant include oxygen, carbon, nitrogen, magnesium, silicon, sulfur, and iron. Lines belonging to some molecules and free radicals, such as OH, NH, CH, and CO, can also be identified in the solar spectrum.
Magnetic fields. Magnetic fields on the sun are measured primarily on the basis of the Zeeman splitting of the absorption lines in the solar spectrum (seeZEEMAN EFFECT). Several types of magnetic fields can be distinguished (see). The sun’s general magnetic field is weak; it reaches a strength of 1 oersted for a given polarity and varies with time. This field is closely connected with the interplanetary magnetic field and the interplanetary field’s sectorial structure. The magnetic fields associated with solar activity can reach strengths of several thousand oersteds in sunspots. The structure of the magnetic fields in active regions is very complicated and magnetic poles of opposite polarity alternate. Localized magnetic fields of hundreds of oersteds are also found outside the sunspots. Magnetic fields penetrate into the chromosphere and into the solar corona.
Magnetohydrodynamic and plasma processes play an important role on the sun. At temperatures of 5000°-10,000°K, the gas is sufficiently ionized and has a high conductivity; because of the vast scale of solar phenomena, the magnitudes of the electromechanical and magnetomechanical interactions are very large (seeCOSMIC MAGNETOHYDRODYNAMICS).
Photosphere. The solar atmosphere consists of the outer, observable layers of the sun. Almost all solar radiation comes from the photosphere, which is the lowest portion of the atmosphere. By making use of the observed radiation flux and the equations for radiative energy transport, radiative equilibrium, and local thermodynamic equilibrium, a theoretical model can be constructed for the distribution of temperature and density with depth in the photosphere.
The photosphere is about 300 km thick, and its mean density is 3 × 10–4 kg/m3. The temperature of the photosphere decreases as the distance from the center of the sun increases; the mean temperature is of the order of 6000°K. The temperature is about 4200°K at the top of the photosphere. The pressure in the photosphere varies from 2 × 104 to 102 N/m2.
The existence of convection in the zone underlying the photosphere is evidenced by the photosphere’s nonuniform luminance and apparent granular structure, which is known as the photo-spheric granulation. The granules are bright patches of a more or less round shape that are visible in photographs of the sun taken in white light. The granules are 150–1,000 km in diameter and last for 5–10 min, although individual granules have been observed for up to 20 min. Sometimes the granules form clusters up to 30,000 km across. The granules are 20–30 percent brighter than the intergranule “lanes.” This difference corresponds to an average difference in temperature of 300°K. Unlike other formations, the photospheric granulation is the same at all heliographic latitudes and is independent of solar activity.
The speeds of the random motions in the photosphere are 1–3 km/sec, according to different determinations. Quasi-periodic radial oscillatory motions have been found in the photosphere. They occur in areas 2,000–3,000 km across and have a period of about 5 min and a velocity amplitude of the order of 500 m/sec. After several periods, the oscillations in a given place die out and may then begin again. Observations have also shown the existence of cells in which motion occurs in a horizontal direction from the center of the cell toward its edges. The speed of the outward flow is about 500 m/sec. The cells, called supergranules, are 30,000–40,000 km across. The supergranules coincide in position with the cells of the chromospheric network. The magnetic field is more intense at the edges of the supergranules. The supergranules presumably reflect the existence of convection cells of the same size at a depth of several thousand kilometers below the surface.
The photosphere was originally assumed to produce only a continuous spectrum, with absorption lines being formed in a reversing layer located above the photosphere. It was subsequently found that spectral lines as well as a continuous spectrum are formed in the photosphere. In order, however, to simplify the mathematics involved in calculating the spectral lines, the concept of the reversing layer is still sometimes used.
Sunspots and faculae. Sunspots and faculae are often observed in the photosphere. Sunspots are dark formations generally consisting of two parts: a darker core, called the umbra, and a surrounding less dark region, the penumbra. The diameter of a sun-spot may be as great as 200,000 km. Some sunspots are surrounded by a bright border. Very small sunspots are called pores. A sunspot’s lifetime varies from several hours to several months.
Even more absorption lines and bands are observed in sunspot spectra than in the spectrum of the photosphere. Sunspot spectra are reminiscent of the spectra of stars of spectral type KO. In what is known as the Evershed effect, the motion of matter in the spots is indicated by Doppler line shifts: gas flows outward from the spot’s center at lower levels and inward at higher levels at speeds reaching 3 × 103 m/sec. Comparison of the intensities of the lines and continuous spectra of sunspots and of the photosphere shows that sunspots have a temperature of 4500°K or less and are 1000°-2000°K cooler than the photosphere. Sunspots consequently appear dark against the background of the photosphere. The luminance of the core of a sunspot is 0.2–0.5 of the luminance of the photosphere, and the luminance of the penumbra is approximately 80 percent of that of the photosphere.
All sunspots have a strong magnetic field. Large sunspots may have fields as intense as 5,000 oersteds. Sunspots usually form groups, which, with respect to magnetic field, may be unipolar, bipolar, or multipolar. The multipolar groups contain many sun-spots of different polarities, with the spots often sharing a common penumbra.
Sunspot groups are always surrounded by faculae, flocculi, and prominences. Solar flares sometimes occur near sunspot groups, and ray and arch structures are observed in the corona above the groups. All these phenomena collectively form active regions on the sun.
The mean annual numbers of observed sunspots and active regions and the mean area occupied by sunspots vary with a period of about 11 years. This period is an average value, and the duration of individual cycles of solar activity varies from 7.5 to 16 years (see). The maximum number of spots simultaneously visible on the surface of the sun is different for different cycles. The maximum in one cycle may be more than twice as great as the maximum in another. Sunspots are generally found in the belts between 5° and 30° heliographic latitude on both sides of the solar equator. The first spots of a cycle occur at high latitudes in these belts, and the last spots appear at low latitudes while spots of the succeeding cycle are appearing at higher latitudes.
Bipolar sunspot groups are observed most often. Such groups consist of two large spots, called the leading and following spots, and often of several smaller spots as well. The leading and following spots of a group are of opposite magnetic polarity. In a given cycle, the polarities of the leading spots are all the same in one hemisphere and are opposite in the other hemisphere.
Sunspots apparently are depressions in the photosphere. The density of the matter in sunspots is less than the density of the matter in the photosphere at the same level.
Faculae are observed in active regions of the sun and are bright photospheric formations visible in white light primarily near the sun’s limb. They usually appear before sunspots and exist for some time after the disappearance of the spots. The area of a fa-cular region is several times greater than the area of the associated group of sunspots. The number of faculae on the solar disk depends on the phase of the solar activity cycle. Faculae have maximum contrast (18 percent) near, but not at, the limb. They are practically invisible in the center of the solar disk, where the contrast is very low. They have a complex, veined structure, and their visibility depends on the wavelength at which observations are conducted. The temperature of faculae is several hundred degrees higher than the temperature of the photosphere. The total radiation from 1 cm2 of a facula is 3–5 percent greater than the figure for the photosphere. Faculae apparently rise somewhat above the photosphere. The lifetime of an average facula is 15 days, but that of individual faculae may reach almost three months.
Chromosphere. The layer of the solar atmosphere that is located above the photosphere is called the chromosphere. Unless special telescopes with narrow-band light filters are used, the chromosphere is visible only during total solar eclipses. At such times, it appears as a rosy arc around the dark limb of the moon during the minutes when the moon completely covers the photosphere. The chromospheric spectrum can also be observed during an eclipse. The spectrum so observed is called the flash spectrum.
The chromosphere at the sun’s limb appears to the observer as an uneven fringe, from which “spikes,” called spicules, emerge. The spicules vary in diameter from 200 to 2,000 km and reach heights of the order of 10,000 km. The spicule plasma moves upward at speeds as high as 30 km/sec. Up to 250,000 spicules may exist simultaneously on the sun.
When observations are performed in monochromatic light, for example, the light of the ionized calcium line at 3,934 A, a bright chromospheric network is visible on the solar disk. This network consists of individual elements. The small elements are 1,000 km in diameter, and the large ones are 2,000–8,000 km in diameter. The large elements are clusters of small elements. The cells of the network are 30,000–40,000 km in diameter. The spicules are presumably formed on the borders of the cells.
When observations are made in the light of the red hydrogen line at 6,563 A, a characteristic vortical structure is seen in the chromosphere near sunspots.
The density of the matter in the chromosphere decreases with increasing distance from the center of the sun. The number of atoms per cm3 varies from 1015 near the photosphere to 109 in the upper chromosphere.
The spectrum of the chromosphere consists of hundreds of emission lines of hydrogen, helium, and metals. The strongest of these spectral lines is the red hydrogen Hα line at 6,563 Å and the H and K. lines of ionized calcium with wavelengths of 3,968 and 3,934 Å. The extent of the chromosphere is different for observations made in different spectral lines. In the strongest chromospheric lines, the chromosphere can be traced out to 14,000 km above the photosphere. Study of chromospheric spectra has led to the conclusion that the temperature is at a minimum in the transition layer between the photosphere and chromosphere, rises to 8000°–10,000°K above the base of the chromosphere, and reaches 15,000°–20,000°K at a height of several thousand km.
Turbulent motions of gas have been found to occur in the chromosphere at speeds of up to 15 × 103 m/sec. In the chromosphere, faculae in active regions are visible in the monochromatic light of the strong chromospheric lines as the bright formations usually called flocculi. The dark formations called filaments are easily seen in the Hα line. At the edge of the solar disk, the filaments extend beyond the disk and are observed against the background of the sky as bright prominences.
Filaments and prominences are most common in four zones that are symmetrically arranged relative to the solar equator: two polar zones to the north of +40° latitude and to the south of + 40° and two low-latitude zones at about ± 30° at the beginning of a cycle of solar activity and at about ± 17° at the end of a cycle. The filaments and prominences of the low-latitude zones exhibit a pronounced 11-year cycle, and their maximum coincides with the sunspot maximum. For the high-latitude prominences, the dependence on the phases of the solar activity cycle is less marked, and the maximum is reached two years after the sunspot maximum.
Filaments that represent quiescent prominences may be as long as one solar radius and may exist for several rotations of the sun. The mean height of prominences above the surface of the sun is 30,000–50,000 km; their average length is 200,000 km, and their average thickness is 5,000 kilometers.
A. B. Severnyi has found that prominences can be divided into three groups according to the character of the motion involved: electromagnetic prominences, in which motions occur along orderly curved trajectories that are lines of force of the magnetic field; random prominences, in which disorderly, turbulent motions predominate (with speeds of the order of 10 km/sec); and eruptive prominences, in which the material of an initially quiescent prominence with random motions is suddenly thrown out away from the sun at an increasing speed that reaches 700 km/sec.
The temperature in prominences or filaments is 5000°-10,000°K; the density is close to the mean density of the chromosphere. Filaments that are active, rapidly changing prominences usually undergo marked alteration over a few hours or even minutes. The form and character of the motions in prominences are closely connected with the magnetic field in the chromosphere and corona.
Corona. The corona is the outermost and most tenuous part of the solar atmosphere and extends more than ten solar radii. The corona can be observed during total solar eclipses, when it appears as a pearly white glow about the solar disk covered by the moon. The structural features of the corona stand out clearly; they include helmet streamers, plumes, rays, and polar brushes. Only with the invention of the coronagraph in 1931 did it become possible to observe the corona at times other than during eclipses.
The general shape of the corona varies with the phase of the solar activity cycle: during years of minimum activity, the corona is markedly drawn out along the equator; in years of maximum activity, the corona is almost spherical. In white light, the surface brightness of the solar corona is a million times less than the brightness of the center of the solar disk. The light of the corona results mainly from the scattering of photospheric light by free electrons. Practically all the atoms in the corona are ionized. The density of the ions and free electrons at the base of the corona is 109 particles per cm3.
The corona is heated in much the same way as the chromosphere. The greatest release of energy occurs in the lower portion of the corona. As a result, however, of the high thermal conductivity, the corona is almost isothermal, and the temperature decreases very slowly with increasing height. Energy is carried out of the corona in several ways. In the lower portion of the corona, the transport of energy downward through thermal conduction plays the principal role. The escape of the most rapid particles from the corona leads to a loss of energy. In the outer portions of the corona, much energy is carried away by the solar wind, which is an outflow of coronal gas. The speed of the solar wind increases with increasing distance from the sun: the speed is several km/sec at the surface of the sun and 450 km/sec at the distance of the earth.
The temperature of the corona exceeds 106°K. In active regions, the temperature is even higher and may reach 107°K. Coronal condensations, in which the particle density is increased by a factor of several tens, may form above active regions.
Part of the radiation of the inner corona is composed of emission lines of multiply ionized atoms of iron, calcium, magnesium, carbon, oxygen, sulfur, and other chemical elements. These lines are seen in both the visible and ultraviolet regions of the spectrum. The radio emission of the sun in the meter band and solar X-ray emission originate in the corona. These emissions are enhanced severalfold in active regions.
Calculations have shown that the corona is not in equilibrium with the interplanetary medium. The particle fluxes forming the solar wind stream out from the corona into interplanetary space.
A relatively thin transition layer exists between the chromosphere and corona. In this layer, the temperature rises sharply to the values characteristic of the corona. The conditions in the layer are determined by the energy flux from the corona as a result of thermal conduction. The transition layer is the source of most of the sun’s ultraviolet radiation.
The chromosphere, transition layer, and corona produce all the observed radio emission of the sun. In active regions, the structure of the chromosphere, corona, and transition layer is altered. This alteration, however, has not yet been sufficiently studied.
Flares. In active regions of the chromosphere, sudden and relatively short-lived increases in brightness are observed that are visible simultaneously in many spectral lines. These bright formations may last for several minutes to several hours and are called flares. They are best observed in the light of the Hα hydrogen line. The brightest flares, however, are sometimes visible even in white light.
The spectrum of a flare has several hundred emission lines of various elements, both neutral and ionized. The temperature of the layers of the solar atmosphere that produce light in the chro-mospheric lines is (1–2) × 104°K; in higher layers, the temperature reaches 107°K. The particle density in a flare reaches 10l3–1014 per cm3. The area covered by a flare may be as great as 1015 m2. Flares usually occur near rapidly developing groups of sunspots with magnetic fields of complex configuration. The occurrence of flares is accompanied by the activation of filaments and flocculi and by the ejection of material. The amount of energy released in a flare is very large and may reach 1024–1025joules. The energy of a flare is presumably initially stored in the magnetic field and then rapidly released. The energy release leads to localized heating and the acceleration of protons and electrons. These accelerated particles cause further heating of the gas, emission by the gas in various regions of the electromagnetic spectrum, and the formation of a shock wave.
Flares produce a considerable increase in the ultraviolet emission of the sun and are accompanied by bursts of X rays (sometimes very intense), by radio bursts, and by the ejection of particles with energies of up to 10’° electron volts. Bursts of X rays are sometimes observed without increased brightness in the chromosphere. Some flares, called proton flares, are accompanied by especially intense fluxes of energetic particles (solar cosmic rays). Proton flares pose a danger for astronauts because when the energetic particles collide with the atoms in the hull of a spaceship, bremsstrahlung, X rays, and gamma radiation are produced, sometimes in dangerous amounts.
Effect of solar activity on terrestrial phenomena. The sun is ultimately the source of all forms of energy used by mankind, except atomic energy. These forms include wind energy, the energy of falling water, and the energy released in the combustion of all types of fuel. Solar activity influences processes occurring in the atmosphere, magnetosphere, and biosphere of the earth in many ways (see).
Solar instruments. Observations of the sun are conducted by means of small-and medium-sized refracting telescopes and large reflecting telescopes. In these instruments, most of the optical components are fixed; the solar rays are directed into the horizontally or vertically positioned fixed telescope by means of one (in a siderostat or heliostat) or two (in a coelostat) moving mirrors. In the construction of large solar telescopes, special attention is devoted to high resolution with respect to the solar disk.
The coronagraph is a special type of solar telescope. The image of the sun is eclipsed within a coronagraph by an artificial “moon,” which is a special opaque disk. In a coronagraph, the amount of scattered light is reduced manyfold, and the outermost layers of the solar atmosphere can therefore be observed at times other than during eclipses.
Solar telescopes are often provided with narrow-band light filters that permit observation in the light of a single spectral line. Neutral filters have also been developed with radially variable transmittance; such filters permit observation of the solar corona at a distance of several solar radii.
The larger solar telescopes are usually equipped with large spectrographs with photographic or photoelectric recording of spectra. The spectrograph may also have a magnetograph, which permits study of the Zeeman splitting and polarization of spectral lines and makes possible determination of the magnitudes and directions of the magnetic fields on the sun.
In order to eliminate the blurring effect of the earth’s atmosphere and to investigate solar radiation in such spectral regions as the ultraviolet and infrared, which are absorbed in the earth’s atmosphere, orbiting observatories located outside the atmosphere have been developed. They permit spectra of the sun and of individual formations on its surface to be recorded beyond the earth’s atmosphere.
REFERENCES
Solntse, vol. 1. Edited by G. Kuiper. Moscow, 1957. (Translated from English.)De Jager, C. Stroenie i dinamika atmosfery Solnlsa. Moscow, 1962. (Translated from English.)
Allen, C. W. Astrofizicheskie velichiny. Moscow, 1960. (Translated from English.)
Mustel’, E. R. Zvezdnye atmosfery. Moscow, 1960.
Severnyi, A. B. Fizika Solntsa. Moscow, 1956.
Zirin, H. Solnechnaia atmosfera. Moscow, 1969. (Translated from English.)
Allen, C. W. Astrophysical Quantities. 3rd ed. London, 1973.
E. E. DUBOV
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