6. Perpetual Change

The solar corona - loops, holes and unexpected heat

Total Eclipse Of The Sun

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Still higher, above the chromosphere, is the corona, from the Latin word for “crown”. The corona becomes momentarily visible to the unaided eye when the Sun’s bright disk is blocked out, or eclipsed, by the Moon and it becomes dark during the day. During such a total solar eclipse, the corona is seen at the limb, or apparent edge, of the Sun, against the blackened sky as a faint, shimmering halo of pearl-white light (Fig. 6.1). But be careful if you go watch an eclipse, for the light of the corona is still very hazardous to human eyes and should not be viewed directly.

A total eclipse of the Sun occurs when the Moon passes between the Earth and the Sun, and the Moon’s shadow falls on the Earth. In an incredible cosmic coincidence, the Moon is just the right size and distance to blot out the bright photosphere when properly aligned and viewed from the Earth. In other words, the apparent angular diameters of the Moon and the visible solar disk are almost exactly the same, so that under favorable circumstances the Moon’s shadow can reach the Earth and cut off the light of the photosphere. The Moon then acts just like your thumb with your arm stretched out and pointed at the Sun. At that distance from your eye, a thumb subtends an angle of about 30 arc-minutes, roughly the same as that of the Moon and the Sun.

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A total eclipse does not happen very often. Since the Moon and the Earth move along different orbits whose planes are inclined to each other, the Moon only passes directly between the Earth and the Sun about three times every decade on average. Even then, a total eclipse occurs along a relatively narrow region of the Earth’s surface, where the tip of the Moon’s shadow touches the Earth (Fig. 6.2). At other nearby places on the Earth, the Sun will be partially eclipsed, and at more remote locations you cannot see any eclipse.

The low corona, that is close to the photosphere, shines by visible sunlight scattered by electrons there. This electron-scattered component of the corona’s white light has been named the K corona. It emits a continuous spectrum without absorption lines, and the K comes from the German Kontinum.

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The amount of observed coronal light is proportional to the electron density integrated along the line of sight, so we can use observations of the K corona to infer the density of electrons there (Fig. 6.3). At the base of the corona there are almost a million billion (1015) electrons per cubic meter. These coronal electrons are so tenuous and rarefied that a million, billion cubic meters would only weigh one kilogram. Since protons are 1836 times more massive than electrons, they supply most of the corona’s mass. Since the corona is made from hydrogen atoms, there is one proton for every electron in the hot gas. The mass density in the low corona is about 10-12 kilograms per cubic meter.

The F corona is the more distant component of the corona’s white light. It extends from about two or three solar radii to far beyond the Earth (Fig. 6.3), and is caused by sunlight scattered from solid dust particles in interplanetary space. Unlike the K corona, the spectrum of the F corona includes dark Fraunhofer absorption lines, so the F stands for Fraunhofer. The faint light of the F corona is not polarized, with any preferred direction, but the K corona is, so polarization is another way to distinguish between the two components.

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White-light coronal photographs show that the electrons can be confined within helmet streamers (Fig. 6.4), which are peaked like old-fashioned, spiked helmets once fashionable in Europe. At the base of helmet streamers, electrified matter is densely concentrated within magnetized loops rooted in the photosphere. Further out in the corona, the streamers narrow into long stalks that stretch tens of billions of meters into space. These extensions confine material at temperatures of about two million degrees Kelvin within their elongated magnetic boundaries.

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Near the maximum in the activity cycle, the shape of the corona and the distribution of the Sun’s extended magnetism can be much more complex. The corona then becomes crowded with streamers that can be found close to the Sun’s poles (Fig. 6.5). At times of maximum magnetic activity, the width and radial extension oaf a streamer is smaller and shorter than at activity minimum. Near solar maximum, the global dipolar magnetic field of the Sun swaps its north and south magnetic poles, so a much more volatile corona can exist then.

Natural eclipses of the Sun occur every few years, and can then be seen from only a few, often remote places on the globe. So, scientists decided to make their own artificial eclipses by putting occulting disks in their telescopes to mask the Sun’s face and block out the photosphere’s intense glare. Such instruments are called coronagraphs, since they let us see the corona. The first coronagraph was developed in 1930 by the French astronomer Bernard Lyot, and the corona is now routinely observed with coronagraphs at mountain sites.

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As Lyot realized, coronagraph observations are limited by the bright sky to high-altitude sites where the thin, dust-free air scatters less sunlight. He therefore installed one at the Pic du Midi observatory in the Pyrennes. The higher and clearer the air, the darker the sky, and the better we can detect the faint corona around the miniature moon in the coronagraph. They work best in space, where almost no air is left. Modern solar satellites, such as the Solar and Heliospheric Observatory or SOHO, use coronagraphs to get clear, edge-on views of the corona from outside our atmosphere (Fig. 6.6). Such satellites also use ultraviolet and X-ray telescopes to view the low corona across the face of the Sun, a development that followed the realization of the corona’s million-degree temperature.

The Corona’s Searing Heat

The solar corona defies expectations, for it is hundreds of times hotter than the underlying photosphere, which is closer to the Sun’s energy generating core. The Sun’s temperature rises to more than one million degrees Kelvin just above the photosphere at a temperature of 5,780 degrees Kelvin. Heat simply should not flow outward from a cooler to a hotter region. It violates the laws of thermodynamics, the branch of physics that deals with the movement and transfer of heat. These laws indicate that it is physically impossible to transfer thermal energy by conduction from the underlying photosphere to the much hotter corona. The high temperature of the corona also defies common sense; after all, when you sit farther away from a fire it becomes colder, not hotter. It is as if a cup of coffee was put on a cold table and suddenly began to boil.

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The linkage between the chromosphere and the corona occurs in a very thin transition region, less than 100 thousand meters thick, where both the density and temperature change abruptly (Fig. 6.7). In the transition region, the temperature shoots up from 10,000 to more than a million degrees Kelvin, but the density decreases as the temperature increases in such a way to keep the gas pressure spatially constant. The corona then thins out and slowly cools with increasing distance from the Sun.

The Ultraviolet And X-Ray Sun

For studying the corona and identifying its elusive heating mechanism, scientists look at ultraviolet and X-ray radiation. This is because very hot material – such as that within the corona – emits most of its energy at these wavelengths. Also, the photosphere is too cool to emit intense radiation at these wavelengths, so it appears dark under the hot gas. As a result, the hot corona can be seen all across the Sun’s face, with high spatial and temporal resolution, at ultraviolet and X-ray wavelengths.

Since ultraviolet and X-rays are partially or totally absorbed by the Earth’s atmosphere, they must be observed through telescopes in space. This has been done using a soft X-ray telescope on the Yohkoh spacecraft, and with ultraviolet and extreme ultraviolet telescopes aboard the SOlar and Heliospheric Observatory, or SOHO for short, and the Transition Region and Coronal Explorer, or TRACE.

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Observations at different temperatures can also be used to focus on different layers of the solar atmosphere. As an example, the spectral line of singly ionized helium, He II, at 30.4 nanometers, is thought to be formed at 60,000 degrees Kelvin, and it is therefore used to image structure in the lower part of the transition region, near the chromosphere (Fig. 6.8). The Sun is mottled all over in this ultraviolet perspective, like a cobbled road or a stone beach. Each stone is a continent-sized bubble of hot gas that flashes on and off in about 10 minutes. The whole Sun seems to sparkle in the ultraviolet light emitted by these localized brightening, known as blinkers. About 3,000 of them are seen erupting all over the Sun, including the darkest and quietest places at the solar poles.

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Images at extreme-ultraviolet and X-ray wavelengths have shown that the hottest and densest material in the low corona is concentrated in magnetic loops. Indeed, Yohkoh’s soft X-ray images have demonstrated that the entire corona is stitched together by thin, bright, magnetized loops that shape, mold and constrain the million-degree gas (Fig. 6.9). Wherever the magnetism in these coronal loops is strongest, the coronal gas in them shines brightly at soft X-ray wavelengths (Fig. 6.10, Fig 6.11).

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High-resolution TRACE images at the Fe IX, Fe XII and Fe XV lines, respectively formed at 1.0, 1.5 and 2.0 million degrees, have demonstrated that there is a great deal of fine structure in the coronal loops (Fig. 6.12). They have pointed toward a corona comprised of thin loops that are naturally dynamic and continually evolving. These very thin loops are heated in their legs on a time span of minutes to tens of minutes, after which the heating stops or changes, suggesting the injection of hot material from somewhere near the loop footpoints in the photosphere or below. The erratic changes in the rate of heating forces the loops to continuously change their internal structure.

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The ultraviolet and X-ray emission of the Sun vary significantly over the 11-year cycle of magnetic activity. The ultraviolet emission doubles from activity minimum to maximum, while the X-ray brightness of the corona increases by a factor of 100. At the cycle maximum, when the sunspots and their associated active regions are most numerous, bright coronal loops dominate the X-ray Sun; at activity minimum the bipolar sunspots and their connecting magnetic loops have largely disappeared, and the Sun is much dimmer in X-rays (Fig. 6.13). However, the corona still stays hot at a minimum in its 11-year activity cycle when active regions go away; the million-degree gas is just a lot more rarefied and less intense.

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Not all magnetic fields on the Sun are closed loops. Some of the magnetic fields extend outward, within regions called coronal holes. These extended regions have so little hot material in them that they appear as large dark areas seemingly devoid of radiation at extreme-ultraviolet and X-ray wavelengths (Fig. 6.14). Coronal holes are nearly always present at the Sun’s poles, and are sometimes found at lower solar latitudes. They are routinely detected by instruments aboard the SOHO, TRACE and Yohkoh spacecraft.

Solving The Heating Crisis

The temperature of the million-degree corona is not supposed to be hotter than the cooler photosphere immediately below it. Heat should not emanate from a cold object to a hot one any more than water should flow uphill. For more than half a century, scientists have been trying to identify the elusive heating mechanism that transports energy from the photosphere, or below, out to the corona. We know that sunlight will not do the trick, for the corona is transparent to most of it.

Magnetic waves provide a method of carrying energy into the corona. The ever-changing coronal loops are always being jostled, twisted and stirred around by motions deep down inside the Sun where the magnetism originates. A tension acts to resist the motions and pull the disturbed magnetism back, generating waves that propagate along magnetic fields, somewhat like a vibrating string. They are often called Alfvén waves after Hannes Alfvén who first described them mathematically. He pioneered the study of the interaction of hot gases and magnetic fields, receiving the Nobel Prize in Physics in 1970 for his discoveries in it.

However, once you get energy into an Alfvén wave, it is difficult to get it out. So there may be a problem in depositing enough magnetic-wave energy into the coronal gas to heat it up to the observed temperatures. Like radiation, the Alfvén waves seem to propagate right through the low corona without being noticeably absorbed or dissipated there.

Magnetic loops can heat the corona in another way – by coming together and releasing stored magnetic energy when they make contact in the corona. Internal motions twist and stretch the magnetic fields, slowly building up their energy. When these magnetic fields are pressed together in the corona, they merge, join and self destruct at the place where they touch, releasing their pent-up energy to heat the gas.

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The SOHO spacecraft has provided direct evidence for such a transfer of magnetic energy from the solar photosphere into the low corona. Images of the photosphere’s magnetism, taken with SOHO, reveal tens of thousands of pairs of opposite magnetic polarity, each joined by a magnetic arch that rises above them. They form a complex, tangled web of magnetic fields low in the corona, dubbed the magnetic carpet (Fig. 6.15). The small magnetic loops rise up out of the photosphere and then disappear within hours or days. But they are continuously replenished by the emergence of new magnetic loops, rising up to form new magnetic connections all the time and all over the Sun.

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The idea of powerful energy release during magnetic reconnection is not a new one. It was proposed decades ago to account for sudden, brief, intense explosions on the Sun, called solar flares, that can release energies equivalent to billions of terrestrial nuclear bombs. Converging flows in solar active regions apparently press oppositely-directed field lines together, releasing magnetic energy at the place that they join. The new, reconnected field lines can snap apart, accelerating and hurling energetic particles out into interplanetary space and down into the Sun.

Such bi-directional, collimated and explosive jets of material have been observed in ultraviolet images of the chromosphere outside active regions (Fig. 6.15). The magnetic interaction of coronal magnetic loops, driven together by underlying convective motions, also energizes at least some of the bright “points” found in X-ray images of the Sun. Unlike sunspots and active regions, the X-ray bright “points” are uniformly distributed over the Sun, appearing at the poles and in coronal holes, some almost as large as the Earth. Hundreds and even thousands of them come and go, fluctuating in brightness like small flares, apparently energized by magnetic reconnection.

The fullness of space

The space between the planets, once thought to be a tranquil, empty void, is swarming with hot, charged invisible pieces of the Sun. They expand and flow away from the Sun, forming a perpetual solar wind. The relentless wind was inferred from comet tails, suggested by theoretical considerations, and fully confirmed by direct in situ measurements from spacecraft in the early 1960s.

The reason that space looks so empty is that the Sun’s wind is exceedingly tenuous, even at its origin near the visible Sun. By the time that it reaches the Earth’s orbit, the solar wind has been further diluted by expanding into the increasing volume of space. There are about 5 million electrons and 5 million protons per cubic meter in the solar wind near the Earth. By way of comparison, there are 25 million, billion, billion (2.5 x 1025) molecules in every cubic meter of our transparent air at sea level. The density of the solar wind is so low that if we could go out into space and put our hands on it, we would not be able to feel it.

The Sun’s continuous wind moves at supersonic speeds near the Earth. It travels with two main velocities, like an automobile with one high gear and one low gear. There is a fast component moving at about 750 thousand meters per second, and a slow one with about half that speed.

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At large distances from the Sun, the charged particles in the solar wind drag the Sun’s magnetic fields with them. While one end of the interplanetary magnetic field remains firmly rooted in the photosphere and below, the other end is extended and stretched out by the radial expansion of the solar wind. The Sun’s rotation bends this radial pattern into an interplanetary spiral shape within the plane of the Sun’s equator, coiling the magnetism up like a tightly-wound spring. This spiral pattern has been confirmed by tracking the radio emission of high-energy electrons emitted during solar flares (Fig. 6.17), as well as by spacecraft that have directly measured the interplanetary magnetism in the ecliptic. The ecliptic plane is the plane of the Earth’s orbit around the Sun, and it is nearly coincident with the Sun’s equatorial plane.

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The shape of the interplanetary magnetic field depends on the Sun’s 11-year cycle of magnetic activity. Near activity minimum, the large-scale, global magnetism of the Sun can be described as a simple magnet with north and south poles where large, unipolar coronal holes are located. The northern hole is of one magnetic polarity, or direction, and the southern one of opposite polarity. The negative and positive field lines meet near the solar equator, where a magnetically neutral layer, called a current sheet, is dragged out into space by the out-flowing wind (Fig. 6.18). Near the Sun, the current sheet coincides with a belt of coronal streamers that seem to meander across the star like the seam of a baseball.

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Because the Sun’s magnetic dipole axis is tilted with respect to its rotation axis, spacecraft near Earth detect a warped current sheet (Fig. 6.19). As the Sun rotates, the current sheet wobbles up and down, like the folds in the skirt of a whirling dervish, sweeping regions of opposite magnetic polarity past the Earth.

Where do the Sun's winds come from?

The solar wind has never stopped blowing during the more than three decades that it has been observed with spacecraft. Two winds are always detected – a fast, uniform wind blowing at about 750 thousand meters per second, and a variable, gusty slow wind, moving at about half that speed.

One method of studying the velocities and origin of the solar wind is called radio scintillation. This technique uses two or more radio telescopes to observe very distant, cosmic radio sources that fluctuate, or scintillate, when their radio waves pass through the solar wind. Optically visible stars similarly twinkle when seen through the Earth’s wind-blown atmosphere. The velocity of the solar wind can be inferred from the time it takes the fluctuating radio signal to move between two telescopes. You could similarly determined the speed of a wind-blown cloud by seeing how long it takes for its shadow to move along the ground.

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Yet, until recently, spacecraft measurements were always made near the ecliptic, and the radio scintillation data only hinted at what the flows looked like above the Sun’s poles. Then, in 1994-95 the Ulysses spacecraft made measurements all around the Sun, at a distance comparable to that of the Earth and near a minimum in the Sun’s 11-year activity cycle. Ulysses’ velocity data conclusively proved that a relatively uniform, fast wind pours out at high latitudes near the solar poles, and that a capricious, gusty, slow wind emanates from the Sun’s equatorial regions (Fig. 6.20).

As the winds blow away, they must be replaced by hot gases welling up from somewhere on the Sun. However, since Ulysses never passed closer to the Sun than the Earth does, simultaneous observations with other satellites were required to tell exactly where the winds come from. Fortunately, the Ulysses data were obtained near activity minimum with a particularly simple corona characterized by marked symmetry and stability. There were pronounced coronal holes at the Sun’s north and south poles, and its equator was encircled by coronal streamers.

Comparisons of Ulysses’ high-latitude passes with Yohkoh soft X-ray images showed that coronal holes were then present at the poles of the Sun, as they usually are during activity minimum. Much, if not all, of the high-speed solar wind therefore seems to come from the open magnetic fields in coronal holes, at least during the minimum in the 11-year cycle of magnetic activity. The coronal holes have comparatively weak and open magnetic fields that stretch radially outward with little divergence, providing a fast lane for the electrified wind. The high-speed solar wind squirts out of the nozzle-like coronal holes, like water out of a fire hose.

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The detailed structure a polar coronal hole has only recently been investigated by the SUMER instrument on SOHO. It showed that the high-speed outflow is concentrated at the boundaries of the magnetic network formed by underlying supergranular convection cells (Fig. 6.21). These edges are places where the magnetic fields are concentrated into inverted magnetic funnels that open up into the overlying corona. The strongest high-speed flows apparently gush out of the crack-like edges of the network, like grass or weeds growing in the dirt where paving stones meet. Thus, SOHO has for the first time discovered one of the exact sources of the fastest winds.

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It was once thought that polar plumes might be the main source of the high-speed wind. Since the long, narrow features are the brightest things around in the dark coronal holes, something had to be energizing them (Fig. 6.22). However, careful SOHO measurements indicate that the fast winds are pouring out of the entire coronal hole, with no substantial difference between the narrow plumes and the inter-plume regions. The honeycomb-shaped boundaries of the magnetic network are apparently the main localized source of the high-speed wind in coronal holes.

Comparisons of Ulysses data with coronagraph images pinpointed the equatorial streamers as the birth place of the slow and sporadic wind during the minimum in the 11-year activity cycle. Hot gas is bottled up in the closed coronal loops at the bottom of the helmet streamers. The capricious slow wind can therefore only leak out along elongated, stretched-out streamer stalks. The part that manages to escape seems varies in strength as the result of the effort.

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Coronal loops are found down at the very bottom of streamers, and the expansion of these magnetized loops may provide the energy and mass of the slow component of the solar wind. Sequential soft X-ray images, taken from the Yohkoh spacecraft, have shown that magnetic loops expand out into space, perhaps contributing to the slow wind. When viewed at the Sun’s apparent edge, near the photosphere, coronal loops are seen rising upwards at speeds of some tens of thousands of meters per second (Fig. 6.23).

Getting up to speed

By the time that the solar wind has reached the Earth, it is moving along at supersonic velocities of hundreds of thousands of meters per second. What forces propel it to such high velocities? The expansion of the hot corona is responsible for some of it.

The corona’s expansion will begin slowly near the Sun, where the solar gravity is the strongest, and then continuously accelerate out into space as it breaks away from the Sun, gaining speed with distance and reaching supersonic velocities. Since there is a limit to the amount of energy being pumped into it, the solar wind will eventually reach a limiting asymptotic or terminal velocity, and then cruise along at a roughly constant speed.

The slow wind naturally reaches terminal velocities of a few hundred thousand meters per second as the million-degree corona expands away from the Sun. Additional energy must be deposited in the low corona to give the fast wind an extra boost and double its speed. In technical terms, the fast wind has a velocity and mass flux density that are too high to be explained by heat transport and classical thermal conduction alone.

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You have to look down into the bottom of the corona to investigate the regions where the corona is heated and the solar wind is accelerated. Yohkoh X-ray observations show that the coronal electrons become fully heated at a height of between 0.2 and 0.5 solar radii, or between 140 and 348 million meters, above the photosphere (Fig. 6.24). In addition, the electron temperatures in coronal holes are several hundred thousand degrees cooler than the temperatures of electrons in coronal streamers at the same height. Ulysses measurements of ion temperatures also indicate that the fast polar wind originates in a relatively low-temperature region in the corona.

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Images from SOHO’s Large Angle Spectrometric Coronagraph, or LASCO, suggest that the slow wind takes a long time to get up to speed (Fig. 6.25). Blobs moving along the stalks of helmet streamers have to move out to 20 or 30 solar radii from Sun center to accelerate to speeds of 300 or 400 thousand meters per second. In contrast, the high-speed wind is accelerated relatively close to the Sun. Radio scintillation measurements indicate that the polar wind reaches terminal speeds of 750 thousand meters per second within just 10 solar radii or less (Fig. 6.26). Thus, the fast wind accelerates quickly, like a racing horse breaking away from a starting gate.

Another SOHO coronagraph, known as UVCS for UltraViolet Coronagraph Spectrometer, has measured temperatures and velocities within the source regions of the solar wind from 1.2 to 10 solar radii from Sun center. It has used the Doppler shifts of ultraviolet spectral lines to show that the high-speed solar wind, emerging from coronal holes, accelerates to supersonic velocity within just 2.5 solar radii from Sun center.

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SOHO’S UVCS has additionally demonstrated that heavier particles in polar coronal holes move faster than light particles in coronal holes. Above two solar radii from the Sun’s center, oxygen ions have the higher outflow velocity, approaching 500 thousand meters per second in the holes, while hydrogen moves at about half this speed (Fig. 6.27). In contrast, within equatorial regions where the slow-speed wind begins, the lighter hydrogen moves faster than the oxygen, as one would expect for a gas with thermal equilibrium among different types of particles.

The amazing thing is that the heavier oxygen ions move faster than the lighter hydrogen in coronal holes. That violates common sense. It would be something like watching people jogging around a race track, with heavier adults running much more rapidly than lighter, slimmer youngsters. Something is unexpectedly and preferentially energizing the heavier particles in coronal holes.

Magnetic waves might preferentially accelerate the heavier ions by pumping up their gyrations around the open magnetic fields. More massive ions gyrate with lower frequencies where the magnetic waves are most intense, thereby absorbing more magnetic-wave energy and becoming accelerated to higher speeds.

The ponderous magnetic waves remind us of the waves in a stormy ocean that push heavy logs to shore. The lighter shells twist and spiral about in the pounding surf, rarely reaching the beach. That is why the heaviest debris is sometimes found left on the beach after high tide.

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The magnetic waves probably block high-energy cosmic rays coming into the Sun’s polar regions, repelling them back into outer space (Fig. 6.28). The incoming cosmic rays meet an opposing force, like a swimmer entering the surf on a distant shore or one trying to swim upstream against the current of a powerful river. To put it in more scientific language, the Alfvén waves are very long, so they can resonate with the energetic cosmic rays and oppose their entry into the polar regions.

Termination of the solar wind

All of the planets are immersed in the solar wind that becomes increasingly rarefied as it spreads out into space. It moves past the planets and beyond the most distant comets. Thus, the entire solar system is bathed in the hot gale that blows frsom the Sun, creating a large cavity in interstellar space called the heliosphere.Within the heliosphere, physical conditions are dominated, established, maintained, modified and governed by the magnetized and electrified solar wind.

The size of the heliosphere has been inferred from the twin Voyager spacecraft, cruising far beyond the outermost planets. At the time of writing they are more than 21 years old and approaching 71 AU. Strong shock waves, associated with intense explosions on the Sun, have plowed into the cold interstellar gas at the heliopause, generating a hiss of radio noise detected by the remote Voyagers. Thirteen months before the spacecraft detected the radio hiss, unusually intense eruptions on the Sun generated one of the largest interplanetary disturbances ever observed. From the measured speed of the disturbance, and the time it took to travel to the heliopause and generate the radio signals, the outer edge of our solar system has been located somewhere between 110 and 160 AU, or roughly a hundred times further from the Sun than the Earth. That is where the solar system ends, in a gigantic distant wall of compressed gas that fences off our Sun from the rest of the cosmos.

Solar Winds at Minimum and Maximum Solar Activity

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In more than 17 years of Ulysses operations, the spacecraft has completed three polar orbits, providing a unique perspective from which to study the Sun and its effect on surrounding space.It has followed the complete course of the 11-year solar activity cycle, first during a minimum in the number of sunspots (in 1994 to 1995), then when the sunspot number was at its maximum (2000 to 2001).

Ulysses thus measured the distribution of solar wind velocities at both the minimum and maximum of the solar cycle (Fig. 6.29).During its first polar orbit, at sunspot minimum, it found fast wind over the poles, and slower, variable wind confined near the solar equator.The second polar orbit, performed near solar maximum, revealed variable solar wind at all latitudes including near the poles.During this second polar orbit at sunspot maximum, the Sun’s polar fields disappeared and then reappeared with the opposite polarity or direction.The third orbit over the solar poles is also at an activity minimum (2007 to 2008), but under very different circumstances after a reversal of the magnetic poles of the Sun, permitting studies of the changed magnetic field and its effect on the solar wind, galactic cosmic rays and solar energetic particles.

Source of Solar Wind at Activity Maximum

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At activity maximum, the large polar coronal holes shrink and disappear, smaller coronal holes appear at all solar latitudes, the slow and fast winds seem to emanate from all over the Sun, and the high-speed winds abate. The slow winds seem to be associated with closed magnetic structures, such as active regions and their associated streamers, or small coronal holes in their vicinity, while the fast winds rush out of the interiors of the largest of the smaller coronal holes (Fig. 6.30). Coronal mass ejections briefly provide a noticeable third flow as they pass through interplanetary space near the maximum of the activity cycle.

Outer Boundary of the Solar System

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How far does the Sun’s influence extend, and where does it all end?The relentless solar wind streams out in all directions, rushing past the planets and carving an immense heliosphere in interstellar space (Fig. 6.31).But since the solar wind thins out as it expands into a greater volume, it eventually becomes too dispersed to repel interstellar forces.The winds are no longer dense or powerful enough to withstand the pressure of gas and magnetic fields coursing between the stars. The radius of this celestial standoff distance, in which the pressure of the solar wind falls to a value comparable to the interstellar pressure, has been estimated at about 100 AU, or one hundred times the mean distance between the Earth and the Sun.

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Instruments aboard the twin Voyager 1 and 2 spacecraft, launched in 1977 and now cruising far beyond the outermost planets, have approached this edge of the solar system from different directions, Voyager 1 moving in the northern hemisphere of the heliosphere and Voyager 2 in the southern hemisphere (Fig. 6.32).Voyager 1 crossed the termination shock of the supersonic flow of the solar wind on 16 December 2004 at a distance of 94 times the mean distance between the Earth and the Sun, or at 94 AU from the Sun.Voyager 2 crossed the termination shock on 30 August 2007 at a distance of 84 AU from the Sun. It appears that there is a significant north/south asymmetry in the heliosphere, likely due to the direction of the local interstellar magnetic field.

Both Voyager 1 and 2 have therefore now crossed into the vast, turbulent heliosheath, the region where the interstellar gas and solar wind interact, due to the reflection and deflection of the solar-wind ions by the magnetized wind beyond the heliosheath.In technical terms, the solar-wind ions in the heliosheath are deflected by magnetosonic waves reflecting off of the heliopause, causing the ions to flow parallel to the termination shock toward the heliotail.

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The motion of the interstellar gas, with its own wind, compresses the heliosphere on one side, producing a teardrop-like, non-spherical shape with an extended tail.A bow shock is formed when the interstellar wind first encounters the heliosphere; just as a bow shock is created when the solar wind strikes the Earth’s magnetosphere. And the graceful arc of a bow shock, created by an interstellar wind, has been detected around the young star LL Orionis (Fig. 6.33).