The story of what we know of the Sun is also the history of solar instruments: the telescope, spectroscope, and many others. Without them we would know little more of the Sun than its position in the sky. Indeed, man's understanding of the Sun has never proceeded regularly, but always in surges, following new instrumental break-throughs. Each new look, from Galileo's simple telescope to Skylab's powerful array, has brought a new and often surprising view of what was once thought to be a simple sphere of fire. Each new look has shown the Sun to be more active and more complex. In many cases, new ways of looking have revealed new and previously hidden faces of the Sun, for the Sun's appearance changes dramatically from layer to layer, as though it wore mask on mask on mask.
THE SUN IN WHITE LIGHT
The mask most commonly seen, and the first to be studied in detail, is the Sun's white light face: the layer in which sunspots are seen, and which astronomy calls the sphere of light or photosphere. It is the white, bright surface we see with the naked eye-the luminous, 6000 K surface that radiates most of the Sun's light and heat to space. To most of us it is the Sun.
The photosphere is a very thin layer of the Sun-no more than about 500 km thick, or less than 0.1 percent of the solar radius. We can see no deeper than this thin shell; although gaseous and diffuse, it becomes completely opaque in a relatively short distance, and that is why the edge of the Sun appears sharply defined, and why early astronomers were led to suppose it could be solid or liquid. The photosphere was long thought to be perfectly spherical and without blemish, but in the early years of the 17th century, with the invention of the astronomical telescope, man's eyes were shown the first details of its intricate and changing surface.
The first and most remarkable features found were sunspots. Dark features on the Sun have been seen and recorded for at least 2000 years. Records from China and the Orient report, on the average, about 5 or 10 sightings per century, which are taken now as evidence of unusually large spots or sunspot groups, because smaller features could not be seen. It is not difficult to see large sunspots with the naked eye if one looks quickly when the.....
 ....brilliance of the Sun is dulled in setting or in rising, or when it is reddened by heavy smoke or haze. But few, if any, of the early records seem to have associated the spots with the Sun itself, probably because tile dark spots were so poorly and intermittently seen. They could be mistaken for dark features floating in our own atmosphere, or planets between Earth and the Sun.
The first telescopes, introduced by Galileo and several other European astronomers in about 1610, did not properly discover sunspots, but rather gave the first clear looks at them. They revealed that the spots were more than dots. Galileo and others demonstrated in careful, published drawings that the spots consist of a dark, central region called the "umbra," meaning "shadow," surrounded by a lighter rim called the "penumbra," or "partial shadow."
 Better telescopes showed sunspots in better detail. The sunspots are usually irregular in shape, some large and some small, and they most often appear in groups, frequently with two principal spots per group. Penumbras are made up of filaments that stream outward from the central umbra in radial fashion. Individual sunspots persist from several days to as long as several months, generally increasing in size and complexity before finally fading away.
Once the spots were shown to be a part of the Sun, scientists used them to demonstrate that the Sun rotated, by watching and measuring their apparent drift across the face of the Sun from day to day.
The Sun rotates in about 27 days. It does not rotate like a rigid ball; the Sun's atmosphere spins faster at the equator than near its poles, introducing a kind of atmospheric slippage or shear known as "differential rotation."
But what were these dark spots in the photosphere?
Most early astronomers believed them to be clouds drifting over an undisturbed and still comfortably perfect solar surface. In 1801, Sir William Herschel, discoverer of the planet Uranus and the leading astronomer of his day, published a different view. Sunspots, Herschel said, were not clouds but holes in clouds. The white photosphere, in his opinion, was an incandescent cloud layer that covered a darker, cooler layer below; sunspots offered occasional views of this lower surface, which, Herschel felt,
Other explanations of the 19th century attributed sunspots to storms, to bubbles on a liquid Sun, splash marks of meteor impacts or of debris thrown out by solar eruptions, mountains poking through the clouds, or solid slag on a white-hot, molten surface.
These fanciful theories of but a century ago are of interest because they reveal the limitations of astronomy without physics: of pictures and visual observations without the benefit of accompanying measurements and deductions of physical quantities, such as temperature, density, velocity, and chemical composition.
The use of physical methods in astronomy, now called astrophysics, began in earnest in the latter.....
 .....part of the 19th century, with the use of the spectroscope. This device, which we will soon describe, was the second of the instrumental breakthroughs that released an avalanche of new knowledge of the Sun.
The first studies of the spectrum of light from sunspots, using a spectroscope, showed that they were cooler by about 30 percent than the surrounding photosphere. A typical temperature for the darkest and coolest part of a sunspot is about 4300 K, which is still hotter than an acetylene welding flame. The photosphere, in contrast, is about the temperature of an electric welding arc.
The spectroscope also revealed that material in sunspots flowed generally upward and outward from the umbra, along the penumbral filaments, suggesting that sunspots were seats of strong, circulatory motion, perhaps like terrestrial tornadoes. If sunspots were solar storms, as many then supposed, they were gigantic and persistent by terrestrial standards. Many of the spots observed were as large as Earth, and some were 10 times larger-the size of Jupiter. They also were longer lasting than terrestrial storms.
An adaptation of the spectroscope by George Hale of the Mt. Wilson Observatory in the early years of this century provided the real key to the understanding of sunspots. Hale found that within sunspots were magnetic fields of such enormous strengths that they dominated all material in their vicinity. This provided an explanation for almost all the observed features of sunspots: their cooler temperatures (magnetic fields inhibited the normal flow of heat toward the surface), the arched penumbral filaments (matter arranged on magnetic lines of force), the frequent occurrence of spots in pairs (of opposite magnetic polarity), and their long persistence.
The Sunspot Cycle
Sunspots, to many, are best known in connection with the sunspot cycle: the regular pattern in the number of spots that appear on the Sun from year to year.
Continuous observation for more than two centuries has established that sunspots reach a maximum number every 11.2 years, on the average. Between times of maxima, their number falls to a well-defined minimum. At the time of maximum in the sunspot cycle we may find 100 or more spots on the visible hemisphere of the Sun at one time;...
 ...at the time of sunspot minimum very few are seen, and for periods as long as a month none can be found.
The last maximum of the sunspot cycle was in 1968-69, and the one before that, which was particularly strong, was in 1957-58. The next should come in early 1980. The last minimum of sunspots was in 1976. Skylab took its observations of the Sun in the declining phase of the sunspot cycle, as the number of spots decreased toward the 1976 minimum.
To the astronomer, this long, clear record of sunspot numbers is evidence that the magnetic disturbances responsible for spots are driven by some cyclic mechanism, probably involving the forces of the differential rotation of the Sun, at and below the level of the photosphere. The search for a complete description of the cause of sunspots and the solar cycle has been one of the long and continuing efforts of theoretical astrophysics.
Other interest in the sunspot cycle springs from its frequent association, in fact and fiction, with cyclic changes on Earth.
Soon after the sunspot cycle was discovered, more than a century ago, a strong relationship was noted between the number of spots on the Sun and the number of auroras that were seen in the northern (and later the southern) skies. A similar relationship was soon established between the number of sunspots and the state of Earth's magnetic field: When the Sun showed more spots, Earth's magnetic field was more frequently disturbed by violent "magnetic storms." In addition, it has been known for many years that the state of Earth's ionosphere-the ionized layer of the high atmosphere that makes possible long-distance radio communication-is affected very critically by solar activity, in step with the number of sunspots. In each of these cases, however, it is not the spots themselves that bring about the terrestrial changes, but other, less frequent and more dynamic events such as solar flares and eruptions in the chromosphere and corona. These events occur more often when the Sun is more spotted and active. Because sunspots are the most easily observed of all solar features, they have traditionally served as the storm warnings of the solar system.
Can these warnings also serve to predict more immediate effects on Earth, such as changes in our weather and climate?
This important question, so frequently asked, is not easily answered. Change in solar output is only one of a number of possible causes of change in the lower atmosphere of Earth. The Sun is a likely cause of long-term climate change as measured in time scales of millenia, centuries, and possibly decades. It is also a possible cause of shorter term climate change. It is probably not a dominant cause of day-to-day or local weather patterns. These are at present only theories, because measurements and physical understanding in this area of solar and atmospheric physics are very limited.
Before more satisfactory answers can be given, we need facts in two specific areas of research. The first is the still unanswered question of what and how much change occurs in the various outputs of the Sun as the sunspot number waxes and wanes. This includes not only its various components of heat and light but also the flow of solar particles in the space between the Sun and Earth. The second need is for a more complete understanding of worldwide circulation, weather, radiation, and climate. Atmospheric physicists are now able to model the complexities of global weather and climate with advanced computers. On the solar side, future measurements from space should add much to the scattered data on solar output that now exist, especially from instruments or programs that continue through a complete solar cycle of 11 years or more.
The photosphere, in spite of its white-hot temperature and intense magnetic fields, is a relatively quiet layer of the Sun, as compared with the more turbulent layers above it. But closer looks at the white light surface, with good telescopes under good observing conditions, reveal an overall pattern of change and motion. The photosphere, under high magnification, appears completely covered with an irregular pattern of cells called granules, which look something like the polygonal cells of a honeycomb.
To scientists who study fluid and gaseous motions, these cells betray the presence of convection, as in the case of turbulent, billowing clouds or bubbles in a boiling pot. Convection in the photosphere carries hot gas upward in globs or cells; these granules radiate some of their heat  energy and then sink back down. Individual cells move up and down at relatively slow speeds (about 18 km/min and last about 8 min, on the average, before breaking up; the solar granulation pattern is thus constantly changing. Although individual granules are hard to detect at our distance from the Sun, they are about 1800 km in diameter, nearly half the size of the continental United States. As we shall see, the up and down motions of the photospheric granulation become more evident and amplified in the solar layers above the photosphere.
Skylab observed the photosphere in white light with several instruments and kept its own photographic record of sunspots, chiefly for comparison with observations in other wavelengths and with solar observations made on the ground during the Skylab mission.
Through the Spectroscope: the Chromosphere, Corona, and Prominences
Above the photosphere and sunspots lie two extensive regions of the Sun that we can call its outer atmosphere-the chromosphere and corona. The chromosphere, or "color sphere," is a layer several times thicker than the photosphere. It is still a thin layer, extending above the photosphere to a minimum height of about 2000 km. It ends in a ragged, irregular boundary with occasional spiked extensions, called spicules, that shoot up an additional 5000 to 15 000 km into the vaporous "crown of the Sun" or corona. Spicules rise and fall like choppy waves on a stormy sea, in periods of 5 to 10 min. The top of the chromosphere, like the top of the ocean, is thus in constant and dynamic change.
The density in these high layers of the Sun is but a fraction of that in the photosphere, which is itself nearly a vacuum like that of the uppermost atmosphere of Earth. Because they are so diffuse and so much fainter than the photosphere, the chromosphere and corona cannot be seen with the naked eye or simple telescope. They can be detected, momentarily, at eclipse, when the Sun blots out the photosphere, or seen in more limited fashion by carefully excluding the glare of the solar disk and looking at the extreme limb of the Sun with instruments designed for the purpose. A most useful tool for observing these outer regions-and their real discoverer-is the spectroscope, which.....
.....can be used to sort out and isolate solar rays that come from specific regions of the Sun.
Physical Conditions in the Sun's Outer Atmosphere
One of the early surprises of solar physics, and for many years one of its perplexing problems, was the elevated temperature of the Sun's outer atmosphere. We generally expect temperature to decrease with distance from a radiating surface. The temperature drops, for instance, as one moves away from a radiator, or the surface of a stove. We would therefore expect the temperature of the solar atmosphere to drop steadily with distance above the radiating photosphere. Instead we find that temperatures in the chromosphere and corona rise sharply and dramatically.
The explanation of this apparent anomaly rests on two facts: the photosphere is not a simple radiator of energy, and the conditions in the upper solar atmosphere are not controlled by radiation alone. Energy of another kind is present, and important, in the chromosphere and the corona-the energy of motion in granules, waves, and spicules. Astronomers now believe that some of the photosphere's energy is transmitted by these means to the thin upper atmosphere of the Sun, and there deposited in the form of heat. Meantime, actual radiation from the photosphere seems to pass unhindered and unabsorbed through the turbulent....
 ....chromosphere and corona. The two mechanisms produce two different effects: (1) radiation pours out the energy that we see and feel on Earth and (2) convection and mechanical motion create and maintain the hot, active, and nearly invisible outer atmosphere of the Sun itself.
Most of the kinetic energy of the granules is deposited in the chromosphere. There, after a slight initial fall, the solar temperature takes its upward turn. While the density drops rapidly (from about 1016 to about 109 particles/cm3 ) the temperature rises from a minimum of about 4300 to about 500 000 K.
At total eclipse, the chromosphere appears to the naked eye as a momentary bright pink band at the very rim of the Sun. It has been seen in this way in fleeting form since at least the early 18th century; but it took the spectroscope to identify the chromosphere as part of the Sun, and to reveal its peculiar physical properties.
The spectroscope, in simplest form, is a glass prism that separates ordinary white light into its....
....component wavelengths or colors; red, orange, yellow, green, blue, and violet. An added eyepiece or projection lens allows the user to see the spectrum clearly and sharply. When a camera replaces the eyepiece, the instrument is called a "spectrograph," which records images of the spectrum called "spectrograms." Other instruments, called "spectrometers," measure the spectrum quantitatively to produce numerical values for the intensity at each wavelength.
Sir Isaac Newton first explained how white light is separated into a rainbow of colors by a prism. Not long after, in 1802, another English physicist, William Wollaston, tried more elaborate experiments with the spectroscope and sunlight. To determine whether sunlight was made up of finer divisions than the broad blur of colors, he inserted a narrow slit between the incoming sunlight and the prism so that images of the slit would show up in the colored spectrum. Wollaston, who was partially blind, found that the normal rainbow of color of the solar spectrum was then crossed by seven dark lines-two in the red region, three in the yellow-green, one in the blue, and one in the violet.
In the following decade Joseph von Fraunhofer, in Munich, performed the same experiment more carefully. He found more than 500 of these dark divisions, which have come to be known as Fraunhofer lines.
Later improvements in instruments revealed more and more lines in varying strengths and.....
...widths. Prisms, by and large, have been replaced by ruled gratings in modern spectographs to achieve high spectral resolution. Today we know that lines in the solar spectrum exist in almost countless number, and that Wollaston and Fraunhofer detected only the darkest and most distinct.
Later 19th century physicists established the remarkable significance of the Fraunhofer lines: each is the unique signature of a distinct chemical element at a distinct stage of ionization in the solar atmosphere. The line seen by Wollaston in the violet part of the spectrum is due to the presence of singly-ionized calcium, for instance; and the strongest of the lines he saw in the yellow is due to neutral, or nonionized, sodium. The line in the blue and one of those in the red are marks of solar hydrogen.
It was later shown, by laboratory experiment,  that the darkness or intensity of each Fraunhofer line is an indication of the amount of the element present, and its width a measure of the temperature and pressure at its source. Slight shifts in the positions of the Fraunhofer lines reveal the presence of motions on the Sun, for velocity of the source changes the wavelength of radiation, by the common phenomenon of doppler shift. It was from these spectrum shifts that sunspot and granulation motions were detected and studied. Other characteristics of the lines can be interpreted to reveal other physical quantities. Magnetic fields, for instance, cause certain spectral lines to blur and split. The work of deciphering the message of the solar spectrum continues today, as laboratory and theoretical studies uncover new methods of interpreting information hidden in the spectrum of the Sun.
Spectra of Photosphere and Chromosphere
Since each layer of the solar atmosphere is characterized by a distinct range of temperature and pressure, each displays its own distinctive spectrum. The spectrum of the photosphere, as we have seen, contains thousands of fine, dark Fraunhofer lines on a bright background, or continuum, of color. The spectrum of the chromosphere, on the other hand, is a reversal of the photospheric spectrum-the continuum of colors is much weaker, and at the wavelengths where dark Fraunhofer lines appeared in photospheric spectra there now appear lines that are brighter than the background. These bright lines are known as emission lines. The brightest of the chromospheric lines in the visible light spectrum corresponds with one of the dark lines seen by Wollaston in the red part of the spectrum, due to hydrogen. Under chromospheric conditions of low density and high temperature, hydrogen emits red light in this narrow spectral region, which is known to astronomers as the line. It is this red emission line that gives the chromosphere its rosy tint and its name.
Spectrum of the Corona
Above the chromosphere, in the solar corona, the spectrum is altogether different; a very weak continuum of color is crossed by only a few lines. These are faint emission lines, only slightly brighter than the background continuum, with positions (i.e., wavelengths) that do not match any of the Fraunhofer lines.
The anomalous spectrum of the corona was first observed on August 7, 1869, at a solar eclipse in Iowa. With a prism spectroscope in a makeshift observatory in Burlington, the American astronomer Charles Young found a strange green emission line in the outermost atmosphere of the Sun. William Harkness, watching the same eclipse with a U.S. Naval Observatory party in Des Moines, observed the same feature. The coronal green line, and others like it, which were discovered at subsequent eclipses, defied explanation for many years.
Spectral lines of the photosphere and chromosphere were associated, one by one, with chemical elements found on Earth, but the lines from the coronal spectrum matched nothing that could be reproduced in the laboratory. For about 70 years they were ascribed to a hypothetical new element, coronium. In 1940 the emission lines of the corona were finally recognized as features of common metals at extremely high temperature and very low densities. This discovery gave the first clear indications of the true temperatures of the Sun's corona- 106 K and more. The green coronal line was shown to be due to atoms of iron, which at these million-degree temperatures and vacuum densities lose half their electrons in successive ionization stages. Other coronal lines in the visible spectrum correspond, we now know, to iron, calcium, nickel, and other heavy elements, which are even more highly ionized, revealing the existence of local regions in the corona that are even hotter.
Spectrum of the Chromosphere
Charles Young was also the first to see the reversed spectrum of the chromosphere. In 1870, the year following his Iowa discovery, Young took the same spectrograph to Spain to observe the next total solar eclipse, which would occur 3 days before Christmas. For this eclipse he kept the slit of the spectrograph tangent to the extreme edge of the Sun so that, at the instant of totality, he might see the unknown spectrum of the chromosphere. During the partial phase, as the Moon covered the last bit of the Sun, he watched the colored continuum and dark lines of the photospheric spectrum fade. Then, at the instant when the  photosphere was completely covered, Young saw the Fraunhofer lines flash from dark to bright, revealing the hidden secret of the chromosphere. It lasted but 2 seconds.
Young's description of this important discovery is a classic in the literature of the Sun:
On his return to America, Young continued his studies of the spectrum of the chromosphere, first in the observatory of Dartmouth College and later in the clear high-altitude air of the Continental Divide.
At the summit of the newly laid Union Pacific Railway, near present-day Tie Siding, Wyo., Young spent the summer of 1872 close to the eyepiece of his telescope, waiting for the sporadic moments when he could best observe the extreme edge of the Sun. During this time he identified and studied over 200 emission lines in the chromosphere. He noted, among other things, that the strength of the chromospheric lines varies from day to day and from position to position around the Sun-an indication that the chromosphere is an active and constantly changing region.
Another dedicated observer who spent many years at the eyepiece of a spectroscope studying the edge of the Sun was Father Angelo Secchi of Rome. Secchi was particularly interested in the ragged upper boundary of the chromosphere, which he described in 1877 as "a burning prairie." In common with other astronomers of his day and since, Secchi was also interested in the larger, more spectacular features at the solar limb called prominences.
Like the chromosphere, prominences were first seen at eclipse.
During a total solar eclipse, when the Sun is covered, occasional solar prominences appear above the edge of the Moon like red clouds standing well above the top of the chromosphere. They can be seen outside the eclipse with spectroscopes and other special instruments. The smallest prominences are about the size of Earth, and some have been seen that are as large as the Sun itself. Their temperatures, densities, and, therefore, their spectra are like that of the chromosphere, and we now consider prominences as cooler, denser chromospheric material that extends into the hotter, rarer, corona.
Prominences appear in many forms. Some develop into explosive features that lift off the Sun as great eruptions into space. Some take the form of large and delicate loops, in which material appears to rain downward, as though condensing from the corona to fall on the chromosphere below. Others are quiet and lasting features of the solar atmosphere that persist for weeks and months without major change. Their elaborate and varied forms, drawn and cataloged by generations of solar astronomers, are now explained as the result of....
....arched lines of magnetic force that connect strong regions of opposite magnetic polarity in the photosphere and chromosphere. Heavy material in quiet, or quiescent, prominences is supported in the lighter corona by a kind of buoyancy of the magnetic lines of force. Loop-shaped prominences and many of the more active and eruptive forms reveal the presence of cooler material in the corona that is trapped in arched magnetic lines, much as electrons and protons are trapped in the van Allen belts of Earth's magnetic field.
The spectroscope, which has provided us with almost all we know about the Sun, remains the astronomer's most powerful tool. In various forms and modifications, such as the spectrograph, spectroheliograph, and the magnetograph, it is the basic instrument of all modern solar research, and an essential part of every solar observatory. Skylab carried five solar spectrographs, each expressly designed to observe important regions of the solar spectrum normally blocked by Earth's atmosphere.
THE SUN IN SINGLE WAVELENGTHS
Light from the disk of the Sun comes from a variety of levels in the solar atmosphere. Different colors, or more precisely, different wavelengths, come from different regions of temperature and pressure that correspond to different heights. Light at the red end of the visible spectrum originates on.....
 ....the Sun at a slightly lower level in the photosphere than does light at the violet end, and when we look at the disk of the Sun in simple white light-which is a composite of all visible colors-we see a composite, or average height. By looking at the Sun in specific colors we can see specific layers. This can be a powerful tool in probing the solar atmosphere.
The advantage is greatly increased when we choose not broad bands of color, but the specific, narrow wavelength regions of the Fraunhofer lines-when we choose, for example, not red or orange or yellow but one of the strong lines within one of those regions.
Solar radiation in the wavelengths of the dark absorption lines comes from regions of the solar atmosphere that lie above the average level of the photosphere-from the higher parts of the photosphere and the chromosphere above it. Different lines and different parts of individual lines originate in different, distinct layers. As a rule, the weaker lines come from lower regions of the solar atmosphere; the stronger lines sample higher regions. The darkest central cores of the strongest lines in the visible spectrum come from the high chromosphere-the level of the spicules seen at the Sun's limb by Secchi.
In scanning through the visible spectrum with a spectroscope we can scan in depth through the photosphere and the chromosphere, much as with a terrestrial spyglass we can examine closer and more distant scenes by adjusting the focus. The deepest we can see into the Sun with our spectroscopic probe is the low photosphere, in the continuum between the lines at the red end of the spectrum; the highest in the upper chromosphere, sampled in the centers of the strongest lines. We can see much higher than the chromosphere by tuning to the continuum and lines of the very short wavelength ultraviolet and X-ray spectrum.
The astrophysicist uses this spectroscopic "tuning" to probe the physical conditions of the Sun and its specific features. If he directs the slit of his spectroscope on a sunspot, for example, he can examine the different levels of the spot and the atmosphere above it by examining different lines and continuum regions in the sunspot spectrum. From these parts of the spectrum he can reconstruct the various levels that make up the sunspot and determine the temperature, density, motions,
and even magnetic field at each level. At any level, however, he will see in the spectrum only that portion of the sunspot that lies within the area of the spectroscope slit, in general a thin strip on the Sun about 1000 km wide. He can move the slit of the spectroscope across the sunspot to compare different parts of it, but he cannot see the entire spot simultaneously. It is much like looking at a picture on a television screen that has only one horizontal scan line; alterations of bright and dark are apparent, but we see only a slice of the full scene and must imagine the rest of it.
The benefit of seeing an entire image of the Sun in a narrow spectral region was obvious to the earliest users of the spectroscope. This hope was realized in 1889 by an undergraduate at MIT who....
 ....modified a spectrograph at Harvard so that it would produce photographic images of the full Sun in a single spectral line. He did it by building up overlapping images of the spectrograph slit on a moving photographic plate. The young astronomer, then 21, was George Ellery Hale and the instrument was called the spectroheliograph. It produced single-wavelength, or monochromatic, images of the Sun called spectroheliograms.
Another avalanche followed, this time in our understanding of the chromosphere. For the first time man saw not just a slice but the entire seething, changing layer of the Sun that envelops the deceivingly quiet photosphere. A spectroheliogram of the Sun made in the center of the or other strong Fraunhofer lines showed a face of the Sun that was surprising in its complexity and its violent change; it was covered with bright and dark areas that changed from day to day and from hour to hour and mottled with a networklike pattern of spicules.
In the l 930's, spectral filters were developed that were almost as effective as the spectroheliograph in isolating radiation of one wavelength, but that had the advantage of greater speed and better uniformity. These simple devices, which can be added to almost any telescope, have no moving parts. They screen out unwanted wavelengths by optical methods, transmitting only a narrow band, such as thewavelength. When you hold such a filter in front of your eye and look at the Sun, you see the Sun in a single wavelength. In , you see only the hydrogen in a certain temperature region on the Sun. Monochromatic filters are made in a variety of ways. The smallest is easily held in the hand and is very effective for viewing the chromosphere. One variety was used in Skylab's arsenal of solar telescopes for portraying the Sun in .
Plages and Filaments
One of the more important discoveries in early spectroheliograms was the brightening of the chromosphere above sunspots. These brighter areas, which were given the French name "plages," (rhymes with "mirages") appeared before sunspots and lasted after the spots were gone; they were recognized as more fundamental indicators of solar activity than the spots themselves. The plages were evidence of hot, dense areas in the chromosphere that were likely seats of dynamic disruptions called.....
...flares. Plages were later shown to delineate regions of the chromosphere where strong magnetic fields were concentrated.
Spectroheliograms also gave extensive new looks at the solar prominences that previously had been seen only with spectroscopes at the limb of the Sun by scientists such as Secchi in his years of patient observations in Rome. Spectroheliograms made in the light of showed prominences on the disk of the Sun, where they appeared as dark  filaments against the bright chromospheric background. Continuous observations of prominences as they crossed the disk of the Sun revealed their real shape and size, lifetime, and relationship to sunspots and other features of the Sun.
Supergranules and the Chromospheric Network
In similar fashion, spectroheliographs showed spicules not just at the edge of the Sun, projected on the sky, but over the entire solar disk. They revealed that these spikes of chromospheric material were distributed over the Sun in a pattern that resembled a fishnet, and which came to be known as the chromospheric network. The spicule network outlines the boundaries of large cells in the chromosphere, encircling each cell like a fence of stakes.
The fenced-in areas were later shown to be giant circulation patterns, much like photospheric gran-....
 ....ules but on a scale 40 times larger. Like the photospheric granules, the chromospheric cells, called supergranules, cover the entire surface of the Sun. They are typically 30 000 km across and more or less circular in shape, persisting about half a day before being disrupted and replaced. Material in supergranules flows up in the center, spreads out to the edge, where the spicules are, and sinks back down. Supergranules are thus giant convective cells that are related to the smaller scale pattern of convective granulation that lies thousands of kilometers below in the photosphere. The heat and kinetic energy of the photospheric granules is transmitted, probably by wave motions, to the chromosphere, where new circulation patterns of grander scale are set up. Magnetic fields in the chromosphere are concentrated in spicules at the supergranule boundaries, where they are apparently pushed by the outward flow of material.
Monochromatic observations of the solar disk revealed for the first time the truly explosive features of the chromosphere called solar flares. Flares are sudden brightenings of limited regions of the chromosphere. By almost every measure they are the Sun's most catastrophic and energetic events. They have special importance to man, for they are known to produce effects that race through millions of kilometers of interplanetary space to rattle the upper atmosphere and magnetic field of Earth, altering the ionosphere and producing auroras and magnetic storms.
Flares last from a few minutes to more than an hour. They are confined on the Sun to relatively small areas; the largest cover but a few tenths of a percent of the solar surface. Flares occur in the chromosphere and corona, above the sunspot layer, but are related to the strong magnetic fields in spots. We see flares almost exclusively in bright plages that are associated with large and complex sunspot groups. Flares have little effect on the sunspots themselves and are hardly noticeable on the bland, white light face of the Sun. Only one flare was observed in white light before the days of the spectroscope, and only a handful before the days of the spectroheliograph and monochromatic filters, which have provided much of what we know about flares.
When flares are seen near the limb of the Sun in a monochromatic picture, they are sometimes.....
.....accompanied by the spraying out of visible material, although not all flares are associated with ejections of this sort. Some on the disk of the Sun appear more like the sudden ignition of a pool of gasoline.
For many years solar observatories of the world have kept up an almost continual patrol for flares, chiefly by making continuous photographs of the Sun in , with automatic cameras equipped with monochromatic filters. If a flare occurs, and if the sky is clear, the camera will record it in successive frames. Related, but harder to achieve, are efforts to record the spectra of different kinds of flares from start to finish. This requires more constant vigilance and quick responses to direct the spectrograph slit at the specific area of the Sun where a flare is starting. It has many of the aspects of hunting quail or other fast, elusive game that flash up unexpectedly and then are gone.
The vast body of data on solar flares has shown that they occur in an almost continuous range of energy and size, from the largest to events so small that they normally escape detection in visible  wavelengths. Flares occur in step with the sunspot cycle: the more spots on the Sun, the greater the probability of a flare. The largest flares occur most often at the time of maximum sunspot number, although not exclusively then. At times of maximum sunspot number we expect to observe a small flare every hour or two. At the time of sunspot minimum, days and weeks transpire without an observable event.
The continuous patrol for flares has shown that they are associated with distinctive features in the chromospheric magnetic field, in plage regions where the field is most concentrated and intense. Solar flares are now believed to result from the sudden conversion of magnetic energy to the energy of heat, light, and motion. We can appreciate the strength of the chromospheric magnetic fields when we realize that a large flare releases energy equivalent to the explosion of 10 million hydrogen bombs in a few minutes' time. The small fraction of this energy that our tiny Earth intercepts is adequate to produce striking effects in our upper atmosphere, even at our great distance from the Sun. These effects are produced by the ultraviolet, X-ray, and energetic particle emission from flares, and not by an appreciable change in the overall output of the Sun at the time. Flares, although spectacular, add only slight amounts to the total energy received from the Sun-at most about a thousandth of a percent, and that for only a few minutes.
Spectra taken from the ground soon revealed that in and other parts of the visible region we see only a small part of the total radiation of flares. Because flares give evidence of very high temperatures, we expect other strong radiation in the ultraviolet and X-ray regions. Radiation from the Sun in these regions can be observed only from above our atmosphere, making the normally difficult tasks of hunting and recording flares even harder. Early solar rockets succeeded in recording ultraviolet and X-ray emission from the later stages of flares, and the first solar satellites provided early spectral information. But it was left to Skylab to record complete spectra of flares from start to finish with high-wavelength discrimination throughout these important spectral regions. This was the most difficult single challenge of the solar observing program on Skylab.
THE MAGNETIC SUN
Shapes and forms seen in the corona at eclipse, particularly at the rotational poles of the Sun,....
 .....resemble patterns taken by iron filings sprinkled over a bar magnet. These appearances led 1 9th century astronomers to question whether the Sun, like Earth, had some kind of basic magnetic field. Other evidences for possible magnetic forces on the Sun were noted in the arrangement of spicules and other features of the chromosphere above sunspots, as observed with the spectroheliograph.
It was only speculation until June of 1908, when George Hale, the inventor of the spectroheliograph, made the first measurement of magnetic fields on the Sun-in sunspots, as observed at Mt. Wilson solar observatory in southern California. The discovery revolutionized the study of the Sun. In finding strong magnetic fields in sunspots, Hale discovered the force for almost all solar activity and the basis for much of modern solar astrophysics.
The Zeeman Effect
Magnetic fields on an object as distant as the Sun cannot be measured directly. Hale's method in 1908, still used today, measures the field indirectly, through its effects on the spectrum. Hale looked closely at the spectrum of light from a sunspot and compared it with that from an unspotted region of the photosphere. One of the significant differences, which had been noted as early as 1892, was a splitting of certain lines of the sunspot spectrum. For example, the Fraunhofer absorption line from iron at a wavelength of 5941 Å, in the red part of the solar spectrum, appeared in sunspot spectra as a close pair of lines, very slightly separated. The same line, outside a sunspot, appeared single. Certain other lines were blurred, or widened, in sunspots.
Shortly after doubled lines were first seen in sunspots, the Dutch physicist Pieter Zeeman in laboratory study discovered that if a light source is placed in an intense magnetic field, certain of its spectral lines are broken into several components, producing split and blurred lines. Zeeman found that the amount of splitting is a measure of the strength of the magnetic field and the direction of the field relative to the observer's line of sight. The splitting is most pronounced when we look along the lines of magnetic force, rather than at right angles to them. Zeeman also noted that the separate components of a magnetically split line show different kinds of polarization, as detected by a polarizer, such as a modern Polaroid. In the simplest case, certain lines, when viewed along the direction of the magnetic field, were split into two components that showed opposite circular polarizations.
Sunspot Magnetic Fields
Many astronomers, following Zeeman's discovery, must have suspected that the blurring and splitting of spectral lines in sunspots were due to magnetic fields. But proof required careful demonstration, involving a good spectroscope of high resolving power, a good solar telescope at a favorable site, and, perhaps most of all, scientists who were involved in modern physics and its application to astronomy. At Mt. Wilson these conditions flourished, and in 1905 Hale set out to make the test. It took an effort of 3 years-an indication of the difficulty of the measurement- but at the end of that time he had the answer: Sunspots showed unmistakable evidence of strong magnetic fields, the strengths of which were several thousands times greater than the magnetic field of Earth.
Magnetic field strengths are commonly expressed in units named after the great German mathematician and physicist Karl Gauss. The strength of Earth's magnetic field, which orients our compasses, is about 0.3 G near the Equator. Near the terrestrial magnetic poles the field is about 0.7 G. In sunspots, Hale found fields of nearly 3000 G extending over sunspot areas many times larger than Earth itself. Fields of this strength and size are awesome sources of great energy. They completely dominate the distribution and motions of matter in their vicinity, and provide the forces that support prominences and explain flares. They also account for the lower temperatures of sunspots by inhibiting the normal flow of convective heat so that sunspots are cooler (and therefore darker) than the surrounding photosphere.
The discovery of magnetic fields in sunspots suggested a search in other areas of the Sun. Hale's method of detecting magnetic fields-by measuring the splitting of certain spectral lines-was limited by the resolving power of the Mt. Wilson solar spectrograph to very strong fields. Moreover, his method, and others used today, could detect only those magnetic fields or components of fields that  happened to be directed toward or away from Earth. Magnetic fields with lines of force perpendicular to the line of sight could not be measured. Another limitation was even more important. Hale was well aware that the field strengths he measured in sunspots, or elsewhere on the Sun, were averaged over a large area-the area of the Sun included in the long rectangle of the spectrograph slit-an a area at least 1000 km in short dimension. Hale pointed out that the averaging process probably concealed a host of fields of small size and possible large strength.
By 1918 Mt. Wilson astronomers thought they had found a weak, large-scale magnetic field at the rotational poles of the Sun. This was evidence, they felt, of a "general magnetic field" that had been proposed to explain the form of the polar corona. Its strength was at that time measured as 20 G-the lowest figure their instrument could measure.
Measurements With the Babcock Magnetograph
A dramatic improvement in techniques of solar magnetic field measurement was made at Mt. Wilson in 1952, when H. W. and H. D. Babcock developed a far more sensitive magnetograph. The Babcocks' instrument took advantage of the marked difference in polarization of the components of a magnetically split line. It employed better spectrographic equipment and newly developed photoelectric sensors to replace the older photographic plates. The result was a great improvement in sensitivity, so that fields of less than 1 G could be detected. This proved to be another important breakthrough in solar instrumentation and research.
The Babcock magnetograph revealed many classes of magnetic fields on the Sun that had escaped detection in earlier spectrographs. It found strong fields in the chromosphere wherever there were plages and concentrated fields in spicules at the supergranule boundaries. After later modification, it measured magnetic fields in prominences and in the low corona. The magnetograph found that flares most often occur along the neutral lines that separate regions of opposite magnetic polarity in the chromosphere.
Study of accumulated magnetograph data over many years at Mt. Wilson showed that the surface of the Sun is divided into very large regions of common magnetic polarity and weak field strength....
.... that persist for many solar rotations. There are, at any time, as many as 20 of these large-scale, unipolar magnetic regions, with boundaries running more or less along lines of solar longitude, and of alternate magnetic polarity. One of the important discoveries from early space measurements demonstrated that this pattern of large-scale magnetic regions extended far above the Sun into interplanetary space, where it was detected by magnetometers on spacecraft at the orbit of Earth.
In the 1950's and 1960's the Mt. Wilson astronomers found evidence that the polar magnetic field of the Sun, the so-called "general field," reverses magnetic polarity each solar cycle, the change occurring at about the time of the sunspot maximum. A solar cycle later, the polarity switches back again, so that the complete magnetic cycle requires two sunspot cycles or 22 years. Hale had found much earlier that the polarities within sunspot groups seemed to reverse polarity in a similar way. In one cycle, the easternmost spot of a sunspot pair would have one polarity, and the western (following) spot of the same group would have the opposite polarity. In the next solar cycle the polarities would be reversed. The spot group reversals, and the apparent reversal of the polar field, imply that the basic cycle of solar activity is not 11 but 22 years.
Modern measurements of the polar field of the Sun find it to be more nearly 1 G in strength, similar to that of Earth. It is thought to be the accumulation of surface fields that have drifted toward the poles of the Sun under the forces of differential rotation of the solar surface. This accumulated weak field near the poles of the Sun accounts for the form of the polar corona and for the misconception about a "general magnetic field" of the Sun.
Improvements in magnetographs have continued to the present day, with important advances in spatial resolution on the Sun. As expected, as we look at smaller and smaller areas on the Sun, we see evidences of finer and finer magnetic fields, with higher and higher field strengths, confirming Hale's suspicion that the averaging over large surface areas was hiding much of what the magnetic Sun had to show us.
THE SUN ECLIPSED
The parts of the outer atmosphere of the Sun-the corona and chromosphere-were first seen during total eclipses of the Sun, when the brighter photosphere is blocked momentarily by the Moon. Until about a century ago this was the only way the chromosphere could be studied, and until less than 50 years ago, the only way the corona could be seen. As a result, much of what we know today of these outer layers has come from the brief observations made at times of eclipse by many generations of astronomers over many years, often with great dedication, effort, and expense.
Total Solar Eclipses
The path of the Moon about Earth is such that from time to time it passes exactly between Earth and Sun, casting the point of a moving, conical shadow in the direction of Earth. The pointed shadow does not always reach Earth, but when it does, it creates a small circle of night, at most about 300 km across, that races across the globe at supersonic speed.
An observer in the path of this small racing shadow will see the Sun's disk totally blocked by a black Moon for a few minutes. During this short time of totality he can see the red edge of the chromosphere (for a few seconds) and the faint, white corona (for a few minutes) against a dark sky. The sudden and ghostly appearance of the corona at a total eclipse is an impressive and moving sight that has never been completely captured in word or picture.
A total eclipse of the Sun occurs somewhere on Earth in two of every three years, on the average. It offers, for those in its narrow path, a time of totality that lasts at most 7 min. The average is about 2 min. In time, every place on Earth will be crossed by the moving shadow of a total eclipse, but, on the average, it is many hundreds of years between occurrences at the same place.
For many centuries, astronomers have been able to predict when and where solar eclipses will occur, and to hundreds of eclipses in all parts of Earth they have carried their telescopes to wait in the path of totality for a few minutes' chance to study the outer atmosphere of the Sun. The English astronomer Baily, for instance, traveled to Scotland to see his first eclipse in 1836 and to Italy in 1842 for his second. If an astronomer, then or now, made it his business to seek out and observe every total solar eclipse for 50 years, he could expect to accumulate a total of about 3 min's observation of the chromosphere and about 1 hr's observation of the corona, allowing for the chance that some of the eclipses would be blocked from view by clouds.
Until the early 18th century, astronomers used solar eclipses chiefly as tests of their ability to predict the orbits of the Moon and Earth, and to take data to refine their knowledge of the laws of celestial motions. A partial solar eclipse, and the partial phases of a total solar eclipse, were for these early astronomers almost as valuable as the rarer and more difficult to observe moments of totality. Few efforts were made to seek the precise places where a solar eclipse would be total, and fewer still attempted to observe the surroundings of the Sun eclipsed.
This may explain why we find no descriptions of the structure of the corona before 1715, and no descriptions of the red flash of the chromosphere and prominences before 1706. It took another century and more before eclipses were used to observe the Sun as a physical object, following the coming of photography and spectroscopy.
Interest in the chromosphere, prominences, and corona grew rapidly in the 19th century, and eclipse experiments were devised to answer the.....
....crucial questions of the origin of these curious appendages of the Sun. What caused the red prominences? Were they more than clouds? Was the corona a part of the Sun, or was it simply light scattered in Earth's atmosphere, like the glare around street lamps in fog? Could it be a lunar atmosphere, illuminated by the Sun? To arrive at answers astronomers needed more observations- careful observations of the detailed form and spectrum of the outer atmosphere of the Sun. As we have seen, the spectroscope allowed study of the chromosphere and prominences in full Sun, but the corona remained an elusive object, seen only at times of total eclipse.
Drawings of the Corona
Photography was not available to assist in the quest until late in the 1800's; therefore, for many years and many eclipses it was standard practice to...
....make drawings of the corona during the fleeting minutes of totality. This was done in pencil, chalk, watercolor, and oil. It was not an easy task, for the corona is made up of ethereal, wispy forms, and glows by eerie light; to many it appeared to shimmer and vibrate, apparently through optical illusion or emotional excitement; and the time alloted by nature to draw the corona was tantalizing at best, with no previews or second sittings. The sketchers, by and large, were not artists but astronomers, and a tense lot. They had traveled, typically, many thousands of kilometers and spent many months in preparation for a short, dramatic test of steady hand and unfailing eye. Some of the sketches were given to better artists, after the eclipse, for redrawing, which improved the quality and not always the accuracy of the portrait.
As a result, the collection of drawings, made with such difficulty, taught very little about the corona. They demonstrated that it was large and white and structured. But artists standing side by side often produced drawings that bore little resemblance to each other. By 1900, George Hale, then editor of the Astrophysical Journal, remarked in an editorial that
The first photograph of the corona, a daguerreotype, was taken in 1851 at a total solar eclipse in Poland, but really useful pictures were not secured for another 40 years. The corona is not easily photographed, and most of the early eclipse photographers came home with blank or vastly overexposed plates. The brightness of the corona changes from eclipse to eclipse and there is always a great change in brightness from inner to outer corona. Plates varied and there was never a chance for a test exposure. A major problem in those early years was photography itself, which was only coming of age.
There were marked improvements, year by year, in photo emulsions, and by the 1890's they had improved to where the form of the corona could at last be photographed in detail. A leader in this effort was the Lick Observatory in California, which in 1889 began a dedicated program of coronal photography, which ran for over 40 years.
In 1893, John M. Schaeberle of the Lick staff developed a special long-focus telescope (or "camera") for photographing the white light corona at eclipse; it was 12 m long and produced images of the Sun and Moon that were 10 cm in diameter. Schaeberle nicknamed it "Jumbo." The giant camera exposed 36- by 56-cm glass plates that were changed by hand by an operator who stood in the dark inside the camera.
From 1893 until 1932 Schaeberle's jumbo camera was taken to nearly every total solar eclipse, on every continent except Antarctica. Each expedition was a major effort requiring the dedication of a significant part of the Lick Observatory staff and resources for about I year. In 1893, when Schaeberle took the camera to Chile for an eclipse in April, his voyage started in February and ended in July. Each eclipse expedition not thwarted by weather produced about 10 new pictures of the corona, taken in a span of a few minutes. Some were unsuccessful, as when clouds or drenching rain blocked the Sun during the precious minutes of totality, or when problems beset the instrument at some remote site.
Coronal Physics From Eclipse
Eclipse photographs proved that the corona was part of the Sun, and told much more about it. The corona is made up of radiating rays and streamers with shapes like tulip bulbs, the pointed ends of which stretch out for millions of kilometers into interplanetary space. The form of the corona changes from eclipse to eclipse and evolves in a fairly regular pattern coordinated with the sunspot cycle. At sunspot maximum the corona is complex and crowded with streamers. At sunspot minimum its form is simpler, usually with streamers restricted to the solar equator and with prominent plumelike rays at the solar pole.
Observations of the corona at eclipse established its source, composition, density, temperature, and the relationship between coronal shapes and magnetic field lines rooted in the surface of the Sun.
The faint white light of the corona is really light from the photosphere, scattered by free electrons that make up the actual coronal structures that we see. When we look at coronal streamers and plumes we are seeing patterns of electrons in the corona, made visible because, like motes in a sunbeam, they scatter the light that strikes them. Electrons scatter light without changing its color, and because the photosphere is white, the corona appears white, too.
Because electrons are influenced by magnetic fields, it is not surprising that the structures of the corona resemble magnetic lines of force. The arches and loops that we see in the corona are magnetic arches and looped field lines, along which coronal electrons cluster. The polar plumes resemble patterns taken by iron filings near a bar magnet because coronal electrons, like iron filings, follow the lines of force emanating from the surface magnetic fields near the pole of the Sun. Coronal streamers are bulb-shaped near their bases because they form around closed magnetic arches that connect surface magnetic fields of opposite polarity.
But coronal magnetic fields are strong enough to control coronal electrons to only a limited height above the Sun. At a distance of about 106 km, the electrons are free to escape and stream away in what is called the solar wind. Electrons are driven from the Sun, along with protons and other particles, by the heat of expansion of the million degree corona. The steady flow of the solar wind stretches out the ends of coronal streamers and gives them their long, straight extensions into space.
Limitations of Eclipse Observations of the Corona
Almost all of these concepts have come from eclipses. But observations of the corona at eclipse are limited in two respects. First, because eclipse pictures cover such limited periods in the life of the corona, they cannot show its dynamic changes. The most complete series of eclipse photographs is only a snapshot album, with long time gaps of a year and more between one set of pictures and the next. Second, eclipse pictures cannot reveal the three-dimensional shape of coronal forms because photographs show the corona always in two dimen  sions, as though projected against the screen of the sky. They cannot tell whether coronal structures are flat or circular in cross section. They cannot tell whether a coronal streamer points obliquely toward Earth or away from it, a question important in tracking its possible influence on Earth. Nor is it possible to determine how features in the corona are related to prominences and flares and other surface activity.
Efforts were made, from the time of the first drawings and first photographs, to search for changes in the corona that might by chance take place during the short time when the Sun was totally eclipsed. In 1905 the Lick Observatory attempted this by building two additional jumbo cameras, copies of Schaeberle's original, and setting them up along the path of totality. For the eclipse of that year one was built in Labrador, one in Spain, and one in Egypt, beside the Nile at Aswan. The time of totality at each station was about 3 min. Totality occurred first at Labrador, about 90 min later in Spain, and still an hour later in Egypt. The results of this ambitious experiment were typical of other such efforts before and since. The Labrador station did not see the corona because of heavy overcast; the sky was hazy in Spain, though adequate for photos, and worse in Egypt, where a desert sandstorm left the sky bright and the plates poor. A careful comparison of Spanish and Egyptian plates revealed no perceptible changes in the corona over a period of 1 hr.
The same experiment had been tried earlier, and at nearly every subsequent eclipse. Movies of the corona are almost always made at eclipse and are studied frame by frame for possible fast changes during the minutes of totality. In all these studies, there were only a few cases of possible subtle changes, as in the orientation of a polar plume or a slight shift in a streamer, but no clear evidence of the dynamic alteration of coronal form that was suspected to result from dynamic surface events. If changes on the surface of the Sun are to be associated with changes on Earth, we need observations of the media between, which includes the passage of these events through the corona.
The one and only hope of detecting changes in the corona, and in answering the questions of its three-dimensional form, depended upon the invention of an instrument that can see the corona when the Sun is not eclipsed. It was obvious to most that this could only be done on a high mountain, where the sky was dark; many mountaintops were tried. For over 60 years the search went on in many ways and in many lands, for a workable coronagraph, but without the slightest shred of success.
THROUGH THE CORONAGRAPH
In the summer of 1930 the French astronomer Bernard Lyot succeeded in developing the contrivance that so many had desired for so long. His coronagraph was a simple device: a telescope in which the image of the photosphere was blocked by a small metal occulting disk, which served as a miniature Moon. Lenses behind the occulting disk focused the corona in an eyepiece, in a spectrograph, or on film.
These same techniques had been tried by others, long before Bernard Lyot was born. He applied nothing new, except patience, perseverance, and meticulous care. With optical loaffles and apertures he reduced the stray light that was scattered in lenses and other parts of the coronagraph. He carried his instrument through the snow to the top of the highest mountain that he could effectively employ, the Pic du Midi, 2680 m above sea level in the French Pyrenees. Lyot's coronagraph was over 6 m long, but he had designed it so that it could be carried on the back of a man. It was made of collapsible aluminum tubing and weighed but 49 lb.
Early Coronagraph Results
Lyot's first pictures of the corona, to those who expected to see it as at natural eclipse, were surely disappointing. With his first coronagraph, in fact, he could not even see the continuum, or white light corona. But with a spectroscope attached to the coronagraph, Lyot could suppress the bright sky background enough to observe the spectrum of the brightest emission lines coming from the inner corona near the edge of the Sun. From the Pic du Midi he saw the coronal lines that others had glimpsed only at eclipse. He could study them more intensively, and set his camera for long exposures-some for half a day. With these long exposures Lyot discovered new coronal lines in the visible and near-infrared spectrum, and learned many important details and relationships of the emission-line corona.
Lyot's coronagraph could see only the innermost corona-as could other ground-based coronagraphs built on his design. They could observe, with the aid of the spectrograph, the lower parts of  the coronal structures seen at eclipse, a narrow band just above the chromosphere. But these data provided the facts and numbers that led to the first understanding of the physics of the corona as a whole. His coronagraph spectra made possible the identification of the source of the coronal lines. They come from very highly ionized but common elements, ending an embarrassing enigma that had troubled astronomy for 70 years, and revealing a dramatic temperature increase in the outer solar atmosphere.
In time, more sophisticated versions of Lyot's coronagraph were developed, including those that took advantage of the polarization of coronal radiation to detect the white light form. These instruments show the shape and density and some changes in the corona, although they are limited by generally coarse spatial resolution and are generally restricted to the lower corona. Coronagraphs supplement eclipse observations of the corona, but they have never supplanted them because they can never really duplicate the unique observing conditions available at the time of a natural eclipse. No....
...ground-based coronagraph has ever produced a picture of the corona that shows the exquisite form and detail that we see with the naked eye during a total solar eclipse.
The reason for this limitation is the light of the sky. At eclipse we see the corona as at night, against a sky so dark that stars appear. With the coronagraph, we look at the Sun in daylight- against a bright sky background that, near the Sun, is particularly intense. The brightness of the sky near the Sun depends upon the altitude and upon the amount of pollutants present in the air, both natural and manmade. The higher and cleaner the air, the darker the sky and the better we can detect the faint corona with a coronagraph. On the Pic du Midi, Lyot demonstrated that in clean, cold air of the Pyrenees the daytime sky is dark enough that one can detect the corona to a distance of a few minutes of arc above the limb of the Sun, or to about 1.3 solar radii from the center of the disk. Subsequent developments in ground-based coronagraphs have extended this, under the best conditions, to about 2 solar radii. At natural eclipse, by contrast, it is possible to observe the corona with a simple telescope or the naked eye to 5 or 6 solar radii, and sometimes farther.
Coronagraphs above Earth
In the 1940's and 1950's, attempts were made to carry a coronagraph above the tops of mountains, in high-flying aircraft and manned balloons. The attempts were not very successful, in part because the brightness of the high-altitude sky was not well known at that time and because the coronagraphs used were only exploratory models. But they brought about an important improvement in the Lyot coronagraph, which was first developed by John W. Evans for aircraft use. He recognized that the Lyot coronagraph was limited by instrument-scattered light, and that this scattered light must be further reduced for a coronagraph to take advantage of skies darker than those at mountaintop. Evans added an extra external occulting disk mounted on a long stem in front of the first coronagraph lens. The added disk shaded the lens, when properly pointed, from direct sunlight. The external occulting disk reduces the scattered light in the coronagraph very significantly, but it does it at the expense of restricting the view. This type of coronagraph can see only the outer parts of the corona.
 The early efforts with aircraft and manned balloons paved the way for more sophisticated and more successful flights in the 1960's and 1970's, including the use of coronagraphs in rockets and on two of the NASA Orbiting Solar Observatory (OSO) spacecraft, in programs of the U.S. Naval Research Laboratory (NRL). The NRL coronagraph on OS0-7 observed, soon after launch, a dramatic transient disturbance in the corona that seemed to be expelled from the Sun into interplanetary space. It was the first time that such an event had been recognized in the corona in white light, and the culmination of a search that had been pursued for more than a century. The same instrument observed a number of subsequent events, setting the stage for a more intensive and detailed study of coronal transients by a larger and more sophisticated coronagraph on Skylab.
At least 4.5 billion
years, in present state
Mean distance from Earth
1.5 x 108 km
Variation in distance through the
1.39 x106 km (or 109
times the diameter of Earth and 9.75 times the diameter of
Jupiter, the largest planet)
1033 cm3 (or 1.3
million times the volume of Earth)
1030 kg (or 333 000 times the weight of
Magnetic field strengths
Bright, chromospheric network
Emphemeral (unipolar) active
10 to 100 G
Earth (for comparison )
0.7 G at pole
Chemical composition of photosphere
(by weight, in percent):
Density (water= 1):
Mean density of entire Sun
Interior (center of Sun)
Sea level atmosphere of Earth (for
Unit area of surface of Sun
Received at top of Earth's
Surface brightness of the Sun
Compared to full Moon
398 000 times
Compared to inner corona
300 000 times
Compared to daytime sky on Pikes
100 000 times
Compared to daytime sky at Orange,
1 000 times
15 000 000 K
Sunspot umbra (typical)
4300 to 50 000 K
800 000 to 3 000 000
Rotation (as seen from Earth):
Of solar equator
At solar latitude 30
At solar latitude 60
At solar latitude 75
At least 4.5 billion years, in present state
Mean distance from Earth
1.5 x 108 km
Variation in distance through the year
1.39 x106 km (or 109 times the diameter of Earth and 9.75 times the diameter of Jupiter, the largest planet)
1.41 x 1033 cm3 (or 1.3 million times the volume of Earth)
1.99 x 1030 kg (or 333 000 times the weight of Earth)
Magnetic field strengths (typical):
Bright, chromospheric network
Emphemeral (unipolar) active regions
10 to 100 G
Earth (for comparison )
0.7 G at pole
Chemical composition of photosphere (by weight, in percent):
Density (water= 1):
Mean density of entire Sun
Interior (center of Sun)
Sea level atmosphere of Earth (for comparison)
3.83 x 1023 kW
Unit area of surface of Sun
6.29 x 104 kW/m2
Received at top of Earth's atmosphere
Surface brightness of the Sun (photosphere):
Compared to full Moon
398 000 times
Compared to inner corona
300 000 times
Compared to daytime sky on Pikes Peak
100 000 times
Compared to daytime sky at Orange, N.J.
1 000 times
15 000 000 K
Sunspot umbra (typical)
4300 to 50 000 K
800 000 to 3 000 000 K
Rotation (as seen from Earth):
Of solar equator
At solar latitude 30
At solar latitude 60
At solar latitude 75