Earth to Sun distance
    Earth to Sun distance km miles AU Light seconds
    Average Distance 149.6 million 93 million 1 8.32
    Current Distance (May 2017) 150 million 93 million 1 8.3
    Maximum Distance 152.1 million 94.5 million 1.02 8.46
    Minimum Distance 147.1 million 91.40 million 0.98 8.18
    Sun to planets distances
    Distance km miles AU Light minutes
    Sun to Mercury Distance 59.13 million 36.74 million 0.4 3.3
    Sun to Venus Distance 108.2 million 67.2 million 0.72 6
    Sun to Earth Distance 149.6 million 93 million 1 8.32
    Sun to Mars Distance 228.9 million 142.3 million 1.53 12.73
    Sun to Jupiter Distance 779.3 million 484.2 million 5.2 43.3
    Sun to Saturn Distance 1.43 billion 888 million 9.55 79.4
    Sun to Uranus Distance 2.88 billion 1.79 billion 19.2 160
    Sun to Neptune Distance 4.49 billion 2.7 billion 30.07 250
    Sun to Pluto Distance 6.09 billion 3.79 billion 40.72 338.66

    More Facts

    Sun is the closest star to Earth. The Sun is a huge mass of hot, glowing gas. The strong gravitational pull of the Sun holds Earth and the other planets in the solar system in orbit. The Sun’s light and heat influence all of the objects in the solar system and allow life to exist on Earth. The Sun is an average star—its size, age, and temperature fall in about the middle of the ranges of these properties for all stars. Astronomers believe that the Sun is about 4.6 billion years old and will keep shining for about another 7 billion years. For humans, the Sun is beautiful and useful, but also powerful and dangerous. As Earth turns, the Sun rises over the eastern horizon in the morning, passes across the sky during the day, and sets in the west in the evening. This movement of the Sun across the sky marks the passage of time during the day (see Sundial). The Sun’s movement can produce spectacular sunrises and sunsets under the right atmospheric conditions. At night, reflected sunlight makes the Moon and planets bright in the night sky. The Sun provides Earth with vast amounts of energy every day. The oceans and seas store this energy and help keep the temperature of Earth at a level that allows a wide variety of life to exist. Plants use the Sun’s energy to make food, and plants provide food for other organisms. The Sun’s energy also creates wind in Earth’s atmosphere. This wind can be harnessed and used to produce power. While it lights our day and provides energy for life, sunlight can also be harmful to people. Human skin is sensitive to ultraviolet light emitted from the Sun. Earth’s atmosphere blocks much of the harmful light, but sunlight is still strong enough to burn skin under some conditions (see Burn). Sunburn is one of the most important risk factors in the development of skin cancers, which can be fatal. Sunlight is also very harmful to human eyes. A person should never look directly at the Sun, even with sunglasses or during an eclipse. The Sun influences Earth with more than just light. Particles flowing from the Sun can disrupt Earth’s magnetic field, and these disruptions can interfere with electronic communications. The Sun is large and massive compared to the other objects in the solar system. The Sun’s radius (the distance from its center to its surface) is 695,508 km (432,169 mi), 109 times as large as Earth’s radius. If the Sun were hollow, a million Earths could fit inside it. The Sun has a mass of 1.989 × 1027 metric tons. This number is very large. Written out, it would be the digits 1989 followed by 24 zeroes. The Sun is 333,000 times as massive as Earth is. Despite its large mass, the Sun has a lower density, or mass per unit volume, than Earth. The Sun’s average density is only 1.409 g/cu cm (1.188 oz/cu in), which is a quarter of the average density of Earth. The Sun produces an enormous amount of light. It generates 3.83 × 1026 watts of power in the form of light. In comparison, an incandescent lamp emits 60 to 100 watts of power. The temperature of the outer, visible part of the Sun is 5500°C (9900°F). From Earth the Sun looks small, because it is far away. Its average distance from Earth is 150 million km (93 million mi). Light from the Sun takes about eight minutes to reach Earth. This light is still strong enough when it reaches Earth, however, to damage human eyes when viewed directly. The Sun is much closer to Earth than any other star is. The Sun’s nearest stellar neighbor, Proxima Centauri (part of the triple star Alpha Centauri), is 4.3 light-years from our solar system, meaning light from Proxima Centauri takes 4.3 years to reach the Sun. The Sun is so much closer to Earth than all other stars are that the intense light of the Sun keeps us from seeing any other stars during the day. Earth would not have any life on it without the Sun’s energy, which reaches Earth in the form of heat and light. This energy warms our days and illuminates our world. Green plants absorb sunlight and convert it to food, which these plants then use to live and grow. In this process, the plants give off the oxygen that animals breathe. Animals eat these plants for nourishment. All plant and animal life relies on the Sun’s presence. The Sun also provides—directly or indirectly—much of the energy on Earth that people use for fuel (see Solar Energy). Devices called solar cells turn sunlight into electricity. Sunlight can heat a gas or liquid, which can then be circulated through a building to heat the building. The energy stored in fossil fuels originally came from the Sun. Ancient plants used sunlight as fuel to grow. Animals consumed these plants. The plants and animals stored the energy of sunlight in the organic material that composed them. When the ancient plants and animals died and decayed, this organic material was buried and gradually turned into the petroleum, coal, and natural gas people use today. The Sun’s energy produces the winds and the movements of water that people harness to produce electricity (see Wind Energy; Water Power). The Sun heats Earth’s oceans and land, which in turn heat the air and make it circulate in the atmosphere as wind. The Sun fuels Earth’s water cycle, evaporating water from the oceans, seas, and lakes. This water returns to the ground in the form of precipitation, flowing back to the oceans through the ground and in rivers. The energy of water’s motion in rivers can be harnessed with dams. The Sun’s gravitational pull holds the solar system together. The planets, asteroids, comets, and dust that make up our solar system are strongly attracted to the Sun’s huge mass. This gravitational attraction keeps these bodies in orbit around the Sun. The Sun also influences the solar system with its diffuse outer atmosphere, which expands outward in all directions. This expanding atmosphere fills the solar system with a constant flow of tiny, fast, electrically charged particles. This flow is called the solar wind. The region through which the solar wind blows is called the heliosphere. Estimates vary about the extent of the heliosphere, ranging from about 86 to about 100 times the distance between Earth and the Sun. Interstellar winds may give the heliosphere an egg shape. The solar wind spreads out as it leaves the Sun. The point at which the solar wind is so diffuse that it stops having an effect on its surroundings is called the heliopause. The heliopause marks the outer edge of the solar system. Within the heliosphere, the Sun provides most of the heat and light that are present, and the particles in the solar wind interact with the planets and satellites in the solar system. The solar wind causes auroras—displays of colored light—in the atmosphere of Earth’s polar regions. The solar wind also carries remnants of the Sun’s magnetic field, which affect the magnetic fields of the planets and larger satellites. The solar wind pushes the planets’ magnetic fields away from the Sun, turning them into elongated, windsock shapes. For more information, see the Solar Wind section of this article. The Sun is extremely important to Earth and to our solar system, but on the scale of the galaxy and the universe, the Sun is just an average star. It is one of hundreds of billions of stars in our galaxy, the Milky Way, which is just one of more than 100 billion galaxies in the observable universe. The Milky Way Galaxy contains about 400 billion stars. All of these stars, and the gas and dust between them, are rotating about a galactic center. Stars that are farther away from the center move at slower speeds and take longer to go around it. The Sun is located in the outer part of the galaxy, at a distance of 2.6 × 1017 km (1.6 × 1017 mi) from the center. The Sun, which is moving around the center at a velocity of 220 km/s (140 mi/s), takes 250 million years to complete one trip around the center of the galaxy. The Sun has circled the galaxy more than 18 times during its 4.6-billion-year lifetime. A star is a ball of hot, glowing gas that is hot enough and dense enough to trigger nuclear reactions, which fuel the star. In comparing the mass, light production, and size of the Sun to other stars, astronomers find that the Sun is a perfectly ordinary star. It behaves exactly the way they would expect a star of its size to behave. The main difference between the Sun and other stars is that the Sun is much closer to Earth. Most stars have masses similar to that of the Sun. The majority of stars’ masses are between 0.3 to 3.0 times the mass of the Sun. Theoretical calculations indicate that in order to trigger nuclear reactions and to create its own energy—that is, to become a star—a body must have a mass greater than 7 percent of the mass of the Sun. Astronomical bodies that are less massive than this become planets or objects called brown dwarfs. The largest accurately determined stellar mass is of a star called V382 Cygni and is 27 times that of the Sun. The range of brightness among stars is much larger than the range of mass. Astronomers measure the brightness of a star by measuring its magnitude and luminosity. Magnitude allows astronomers to rank how bright, comparatively, different stars appear to humans. Because of the way our eyes detect light, a lamp ten times more luminous than a second lamp will appear less than ten times brighter to human eyes. This discrepancy affects the magnitude scale, as does the tradition of giving brighter stars lower magnitudes. The lower a star’s magnitude, the brighter it is. Stars with negative magnitudes are the brightest of all. Magnitude is given in terms of absolute and apparent values. Absolute magnitude is a measurement of how bright a star would appear if viewed from a set distance away. By convention, this distance is 10 parsecs, or 32.6 light-years. Apparent magnitude measures how bright a star appears from Earth. The Sun’s absolute magnitude is 4.8. The brightest known stars have absolute magnitudes of about -9 (lower magnitudes mean brighter stars), and the dimmest known stars have absolute magnitudes of about 20. The apparent magnitude of the Sun is -26.72. The apparent magnitude of the brightest star in Earth’s night sky, Sirius, is -1.46. The dimmest stars that can be seen from Earth with unaided eyes have apparent magnitudes of about 6. Astronomers also measure a star’s brightness in terms of its luminosity. A star’s absolute luminosity or intrinsic brightness is the total amount of energy radiated by the star per second. Luminosity is often expressed in units of watts. The Sun’s absolute luminosity is 3.86 × 1026 watts. The absolute luminosity of stars ranges from one thousandth of the luminosity of the Sun to 10 million times that of the Sun. Another way of measuring brightness is to measure the amount of light that reaches an observer. This measurement is called apparent brightness or apparent luminosity. Apparent luminosity depends on the absolute luminosity of a star and the distance from the star to the observer. Apparent luminosity becomes smaller as distance from the star to the observer becomes larger. From Earth, the apparent luminosity of the Sun is 10 billion times greater than the apparent luminosity of the next brightest star, Sirius, because the Sun is so much closer to Earth. The radius of the Sun is about average among stars. The radii of most stars fall between 0.2 and 15 times the Sun’s radius, although some giant stars are hundreds of times larger than the Sun. Larger stars usually have larger absolute luminosities. We receive much more energy from the Sun than from other stars, because the Sun is so nearby. The Sun’s proximity also allows scientists to study its face in detail. A modest telescope can resolve solar structures that are 700 km (400 mi) across—about the distance from Boston, Massachusetts, to Washington, D.C. That level of detail is comparable to seeing the features on a coin from 1 km (0.6 mi) away. Other stars are so distant that the details on their surfaces remain unresolved with even the largest telescopes. The Sun is a second-generation star, meaning that some of its material came from former stars. Some stars in our galaxy are nearly as old as the expanding universe, which scientists believe originated in the big bang explosion about 14 billion years ago (see Big Bang Theory). In contrast, the Sun is only 4.6 billion years old. The first stars were composed only of the hydrogen and helium produced in the early universe. These stars are called first-generation stars. Although hydrogen is also the main ingredient of the Sun, it contains heavier elements, such as carbon, nitrogen, and oxygen, as well. These elements formed inside first-generation stars that lived and died before the Sun was born. When these massive, short-lived stars used up their internal fuel, they exploded and ejected the heavier elements into interstellar space. The Sun formed from this material, making it a second-generation star. The Sun and planets in our solar system formed when a rotating cloud of dust and gas in space collapsed, or condensed, due to the gravitational attraction between the particles in the cloud. A nearby supernova explosion may have triggered the collapse, or a random fluctuation in the density of the cloud may have started the process. The Sun formed at the center of the spinning cloud, while the debris that condensed into planets formed a flattened disk revolving around the Sun. When the Sun reached its present size about 4.6 billion years ago, it was hot enough inside to ignite the nuclear reactions that make it glow. The Sun cannot shine forever, because it will eventually use up its present fuel. The nuclear fusion reactions that make the Sun glow (for more information, see the section entitled The Sun’s Energy in this article) depend on the element hydrogen, but the hydrogen in the Sun’s core will eventually run out. Nuclear reactions have converted about 37 percent of the hydrogen originally in the Sun’s core into helium. Astronomers estimate that the Sun’s core will run out of hydrogen in about 7 billion years. The Sun will grow steadily brighter as time goes on and more helium accumulates in its core. Even as the supply of hydrogen dwindles, the Sun’s core must keep producing enough pressure to keep the Sun from collapsing in on itself. The only way it can do this is to increase its temperature. The increase in temperature raises the rate at which nuclear reactions occur and makes the Sun brighter. In 3 billion years, the Sun will be hot enough to boil Earth’s oceans away. Four billion years thereafter, the Sun will have used up all its hydrogen and will balloon into a giant star that engulfs the planet Mercury. At this point in its life, the Sun will be a red giant star. The Sun will then be 2,000 times brighter than it is now, and hot enough to melt Earth’s rocks. At this time the outer solar system will get warmer and more habitable. The icy moons of the giant planets may warm enough to be covered by water instead of ice. When the giant Sun uses up its fuel, it will no longer be able to support the weight of its inner layers, and they will begin to collapse toward the core, eventually producing a small, dense, cool star called a white dwarf. The Sun will then have about the same radius as Earth has, but it will be much denser and more massive than Earth. The Sun will become a white dwarf star about 8 billion years from now. After it becomes a white dwarf, it will cool slowly for billions of years, eventually becoming so cool that it will no longer emit light. The Sun produces an amazing amount of light and heat through nuclear reactions (Nuclear Energy). The process that produces the Sun’s energy is called nuclear fusion. In nuclear fusion, two atoms come together to produce a heavier atom. Fusion reactions release energy and tiny elementary particles. In just one second the Sun emits more energy than humans have used in the last 10,000 years. The Sun has been shining relatively steadily for 4.6 billion years. Until the early 20th century, humans did not know of any process that could explain the energy production of the Sun. Even if a fire, such as those that occur on Earth, were as large as the Sun, the fire would consume the mass of the Sun in a few thousand years. Scientists now know that the Sun is mainly composed of hydrogen, the lightest and most abundant element in the universe. The Sun contains an enormous amount of hydrogen, however, which makes the Sun very massive. All matter inside the Sun is gravitationally attracted to all the other matter in the Sun, and this attraction tends to pull the Sun’s mass together. This inward pull creates high pressures and temperatures inside the Sun. The center is so violent and hot that collisions between atoms break the hydrogen atoms apart into their subatomic ingredients. A hydrogen atom is made up of a nucleus that contains a positively charged proton, and a negatively charged electron that orbits the nucleus. In the Sun, collisions separate the electron from the nucleus, freeing each to move about the solar interior. The positively charged nuclei, or protons, are called ions. A gas in which particles are ionized, or have electric charges, is called plasma. Scientists often consider plasma, such as the material inside the Sun, to be a fourth state of matter—the three more familiar states of matter are gas, liquid, and solid. See also Atom. The separation of hydrogen nuclei from their electrons makes nuclear fusion possible at the Sun’s core, producing the Sun’s light and heat. With their electrons gone, hydrogen nuclei (protons) can be packed much more tightly than complete atoms. At great depths inside the Sun, the pressure of overlying material is enormous, the protons are squeezed tightly together, and the material is very hot and densely concentrated. At the Sun’s center, the temperature is 15.6 million degrees C (28.1 million degrees F), and the density is more than 13 times that of solid lead. This is hot and dense enough to make the nuclei fuse together. Outside the solar core, where the overlying weight and compression are less, the gas is cooler and thinner, and nuclear fusion cannot occur. The nuclear fusion reaction that powers the Sun involves four protons that fuse together to make one nucleus of helium. Two of the original protons become neutrons (electrically neutral particles about the same size as protons). The result is a helium nucleus, containing two protons and two neutrons. The helium nucleus is slightly less massive (by a mere 0.7 percent) than the four protons that combine to make it. The fusion reaction turns the missing mass into energy, and this energy powers the Sun. The relationship between energy and the missing matter was explained in 1905 by German-born American physicist Albert Einstein. The mass loss, m, during the transformation of four protons into one helium nucleus, supplies an energy, E, according to the relation E = mc2, where c is the speed of light. The speed of light is a constant number equal to 3 × 108 m/s (1 × 109 ft/s). Every second, fusion reactions convert about 700 million metric tons of hydrogen into helium within the Sun’s energy-generating core. In doing so, about 5 million metric tons of this matter become energy. This energy leaves the Sun as radiation, and the part of this radiation that constitutes visible light is what makes the Sun shine. The rate of nuclear reactions in the Sun is relatively low, because protons repel each other. This repulsion often prevents them from getting close enough to each other to fuse. Protons push each other away because they have the same electrical charge. The particles must overcome this repulsion in order to fuse together. Only a tiny fraction of the protons inside the Sun are moving fast enough to overpower this repulsive electrical force. The nuclei that are moving fast enough can get very close together, and a force called the strong nuclear force takes over. The strong nuclear force is, as its name implies, very powerful, but only over very short distances. It pulls the nuclei together and holds them together. In this way, nuclear reactions proceed at a relatively slow pace inside the Sun. If the pace were much quicker, the Sun would explode like a giant hydrogen bomb. Four protons do not combine directly to form a helium nucleus, since the protons are constantly moving and are almost never in the same place at the same time. Moreover, the electrical repulsion between four protons is too great to overcome, even if the four protons happen to come together at an appropriate speed at the same time. Instead, the protons come together in a series of steps to form a helium nucleus, and these steps are called the proton-proton chain. In the first step of the proton-proton chain, two exceptionally fast protons meet head on and merge into each other, tunneling through the electrical barrier between them. The two protons combine, with most of their mass forming a deuteron, the nucleus of a heavy form of hydrogen known as deuterium. A deuteron contains one proton and one neutron, so one of the protons must become a neutron in this step. The conversion of a proton to a neutron releases a much smaller particle called a neutrino. There are several types of neutrinos—the type that the proton-proton chain produces is called an electron neutrino. The reaction also creates a positron, a positively charged particle the size of an electron. The symbolic representation of the first step of the proton-proton chain is p + p → 2D + e+ + νe where p represents the protons, 2D represents deuterium, e+ represents the positron, and νe represents the electron neutrino. In the second step of the chain, the deuteron collides with another proton to form a nucleus of light helium, which has two protons and one neutron. Less energy is needed to maintain a light helium nucleus than is needed to maintain a deuteron and a proton separately. The extra energy is released as a photon, or a packet of light energy. In symbolic terms, the second step is 2D + p → 3He + g where 3He is light helium and g represents a photon. In the final step of the proton-proton chain, two light helium nuclei meet and fuse together to form a nucleus of normal heavy helium, which has two protons and two neutrons. This reaction also releases two unattached hydrogen nuclei that return to the solar gas. In symbolic terms, the third step is 3He + 3He → 4He + 2p where 4He represents a normal helium nucleus with two protons and two neutrons. The positron created in the first step of the chain eventually collides with a free electron. The positron and the electron are opposite particles—the positron is the antimatter equivalent of the electron. When the positron and the electron collide, they annihilate each other, releasing energy. The electron and the positron disappear, their mass transformed into two photons: e+ + e- → 2g where e- represents the electron. The net result of the proton-proton chain is the transformation of four hydrogen nuclei into a helium nucleus (with two protons and two neutrons), two neutrinos, and six photons: 4p → 4He + 2νe + 6g. The conversion of two protons into two neutrons in the proton-proton chain produces two tiny, elusive, fast-moving neutral particles called neutrinos. Nuclear reactions in the Sun’s central furnace create prodigious quantities of neutrinos. Every second the Sun releases 2 × 1038 neutrinos, and every second an estimated 70 billion of these solar neutrinos pass through every square centimeter of Earth that is facing the Sun. Neutrinos move at the velocity of light, have no electrical charge, and have so little mass that scientists are not sure that neutrinos have any mass at all. The ghostlike neutrinos therefore travel almost unimpeded through the Sun, Earth, and nearly any amount of matter. Scientists can snag small numbers of neutrinos in massive underground detectors called neutrino telescopes (see Neutrino Astronomy). These telescopes are placed so deep underground that only neutrinos can reach them. Scientists using these telescopes have detected solar neutrinos, confirming that the Sun is indeed powered by nuclear fusion. The number of neutrinos detected by these telescopes, however, is only one-third to one-half of the total number of neutrinos predicted to exist by the theory of solar neutrino production. This discrepancy between the number of detected neutrinos and the number predicted is known as the solar neutrino problem. There are two possible explanations—scientists might not understand exactly how the Sun produces its energy, or they could have an incomplete knowledge of neutrinos. Astronomers are convinced that their models of the Sun are correct and that their predictions for the expected amount of solar neutrinos are therefore correct. Studies of the interior of the Sun substantiate the current models of how the Sun produces its energy, so most scientists agree that the problem lies in their understanding of neutrinos. The theory scientists favor to explain the problem is that neutrinos from the Sun change on their way to Earth. Scientists know of at least three types of neutrinos. Nuclear fusion reactions in the Sun produce a type of neutrino called an electron neutrino. The other two proven types of neutrinos are called muon neutrinos and tau neutrinos. Most neutrino telescopes, especially those devoted to solar research, can only detect electron neutrinos. In the 1990s studies of muon neutrinos (produced by reactions between particles called cosmic rays and Earth’s atmosphere) showed that muon neutrinos might change into tau neutrinos. Research conducted since the late 1990s indicates that electron neutrinos from the Sun may also change into another type of neutrino. This change would mean the electron neutrino detectors miss many of the Sun’s neutrinos. The energy that the Sun produces in its core must travel to the Sun’s surface to make the Sun glow. The mechanisms that transport radiation from the center to the surface of the Sun define the structure and behavior of the layers inside the Sun. Nuclear fusion releases energy deep down inside the Sun’s high-temperature core, which extends from the center to about one-quarter of the radius of the Sun. The layers above the core produce no energy, so the core, which makes up only 1.6 percent of the Sun’s volume, produces all of the Sun’s energy. Energy moves from the core to the rest of the Sun through two spherical shells that surround the core. The inner shell is called the radiative zone, and the outer one is called the convective zone. Radiation and convection are two ways that energy can travel from one place to another (see Heat Transfer). Radiation involves the movement of energy, but not the movement of material. The radiative energy spreads out in all directions and can move between objects that are not connected. Radiation can be absorbed by another substance. In the process of convection, matter moves energy. Convection occurs when a liquid or gas moves into contact with an object at a different temperature. Energy moves from the core of the Sun to the next innermost layer, the radiative zone, through radiation. The radiative zone spans from the outer edge of the core, which is 174,000 km (108,000 mi) from the Sun’s center, to 496,000 km (308,000 mi) from the Sun’s center. The radiation diffuses outward in a haphazard, zigzag pattern. Particles in the radiative zone repeatedly absorb, radiate, and deflect photons of energy. The matter in the radiative zone stays in the same place while the energy moves through it. Because of this continued ricocheting in the radiative zone, about 170,000 years, on average, are required for a photon of energy to work its way outward from the Sun’s core to the bottom of the convective zone. The Sun’s interior cools with increasing distance from the center, as the heat and radiation of the core spread outward into an ever-larger volume. At the base of the convective zone, the temperature is about 2.2 million degrees C (about 4.0 million degrees F). At the boundary of the cooler convective zone, the radiative energy has lost too much intensity and the material is too cool and dense to allow the energy to pass through. The layers of material at the bottom of the convective zone heat up with blocked radiation and become less dense than surrounding material. This heated material then moves up through the convective zone, carrying energy toward the atmosphere of the Sun. When the material reaches the atmosphere—a layer that is much less dense than the convective zone—the energy can radiate into space. The material at the top of the convective zone becomes cooler and therefore denser when it releases its energy, falling back down to the bottom of the zone to pick up more energy. The length of time needed for a particle to pass through the convective zone, from the innermost to the outermost edge, is about ten days. The behavior of the outer, visible layer of the Sun allows scientists to glimpse the structure of the interior of the Sun. The visible part of the Sun is called the photosphere. The photosphere heaves in and out with a rhythmic motion. The material in the photosphere can reach a height of 50 km (30 mi) and speeds of 500 m/s (1,600 ft/s). The time each oscillation takes to go from its highest point to its lowest point and back again is called its period. Each oscillation has a period of about five minutes. The oscillations in the photosphere are actually caused by sound waves from the convective zone. Sound waves, whether on Earth or in the Sun, are waves carried by matter. They travel by compressing matter in their path. Because they rely on matter, sound waves cannot travel through a vacuum, or an area in which no matter is present. Air carries most of the sound we hear on Earth. The hot plasma of the Sun carries sound waves within the Sun. Hot gas churns in the convective zone, producing a noise like that of a jet airplane or a pot of boiling water, but much, much louder. When these sounds strike the photosphere and rebound back down, they disturb the gases there, causing them to rise and fall. The sound waves are trapped inside the Sun and cannot travel through the vacuum of space. Even if they could reach Earth, the Sun’s sounds are too low-pitched for the human ear to hear. A period of five minutes corresponds to 0.003 vibrations per second. The lowest sounds that even a sensitive human ear can hear have a frequency of about 25 vibrations per second. Scientists can “listen” to the Sun’s vibrating notes indirectly by watching the rhythmic motions of the photosphere. Sensitive instruments detect the Sun’s oscillations by recording periodic changes in the wavelength of the Sun’s light. Motion at the solar photosphere changes the wavelength of the light that it emits. When oscillations on the Sun’s photosphere move its material away from Earth, the Sun’s light shifts to longer wavelengths. This shift occurs because each successive wave has farther to travel than the one before it did in order to reach Earth, so the distance between waves (the wavelength) becomes longer. Photospheric oscillations that move material toward Earth make the wavelengths shorter, because each wave has a shorter distance to travel than the one before it did. These changes in wavelength due to motion are called the Doppler effect. Helioseismology is the study of the interior of the Sun. Measuring the speed of sound waves in the Sun helps scientists determine the temperature and composition of the Sun. The speed of sound depends on the temperature and composition of the material through which the sound passes. Helioseismologists exploit this relationship to establish how the Sun’s temperature, density, and composition vary with distance from the center. Experimental measurements of the density and temperature of the Sun’s layers agree almost perfectly with theoretical models. Measurements of the Sun’s core temperature are very close to the predicted value, showing that the predicted number of solar neutrinos should also be correct. Scientists also use oscillations in the photosphere to study the movement of the interior of the Sun. About 10 million separate sounds—each traveling in a different, defined section of the solar interior—combine to produce the oscillations in the photosphere. Scientists can separate all of the different vibrations, trace them back to their origins, and look into the heart of the Sun. Like Earth, the Sun rotates, or spins, around an imaginary line that runs through its center. This line is called the Sun’s axis, and the top and bottom of this line mark the Sun’s north and south poles, in the same way that Earth’s axis marks the North Pole and South Pole on Earth. Earth, the Sun, and the other planets in the solar system all lie on one plane, and the Sun’s north pole and Earth’s North Pole are oriented in roughly the same direction relative to the plane. The Sun’s equator, like Earth’s, is an imaginary line halfway between the north and south poles that runs east and west. Like Earth, the Sun rotates from west to east when viewed from above the north pole, but unlike Earth, different parts of the Sun rotate at different rates. In the photosphere, the areas near the north and south poles of the Sun rotate more slowly than the areas nearer the solar equator. A spot at the Sun’s equator takes 25 days to rotate completely, while a spot 15° from the poles takes 34 days to make a complete rotation. This phenomenon is known as differential rotation. Scientists use helioseismology to measure the Sun’s internal rotation. Sound waves moving in the direction opposite to the rotation of the Sun appear to move more slowly than those moving with the rotation of the Sun. Helioseismologists can pinpoint the origins of fluctuations on the Sun’s surface and compare sound waves that have taken different paths to the surface. Armed with this sensitive indicator, helioseismologists have shown that the differential rotation exhibited by the photosphere persists throughout the convective zone. These differences disappear in the underlying radiative zone, where the rotation speed becomes uniform from pole to pole. At the boundary where the convective and radiative zones meet, the different rotation speeds cause the material in the zones to rub together. Scientists suspect that the forces generated by the two zones moving against each other may create the Sun’s magnetic field. The material in the Sun farther out from the center than the photosphere makes up the Sun’s atmosphere. The atmosphere extends far beyond the disk we see in the sky. Very diffuse solar gases extend all the way to Earth and beyond. The solar atmosphere consists of, from the innermost part outward, the photosphere, the chromosphere, the corona, and the expanding outer layers of the corona that astronomers call the solar wind. The photosphere is the visible part of the Sun. We look right through the chromosphere, the corona, and the solar wind, just as we see through Earth’s atmosphere at night. The chromosphere and corona are visible during total solar eclipses, when the Moon lines up between the Sun and Earth, completely blocking the main disk of the Sun from view. The thin chromosphere becomes visible a few seconds before or after a solar eclipse, creating a narrow pink, rose, or ruby-colored band at the edge of the Sun. For up to eight minutes during an eclipse, the corona is visible to the unaided eye as a faint, shimmering halo of pearl-white light spreading out from the lunar silhouette. Although the light of the chromosphere and corona is not bright enough to be dangerous, and can be viewed safely without filters during the total phase of an eclipse, the partial phases of a solar eclipse are very hazardous to human eyes and can only be viewed indirectly or through special filters. Scientists can study all layers of the Sun’s atmosphere at any time using special instruments. The photosphere is the lowest, densest level of the solar atmosphere. The visible light that reaches Earth from the Sun originates in the photosphere. That light comes from a thin, bright shell about 300 km (about 200 mi) thick, a thickness of less than 0.05 percent of the Sun’s radius. The photosphere has a temperature of 5510°C (9950°F). It is a diffuse, tenuous gas with a pressure that is only a small fraction, 0.0001, of the amount of pressure in Earth’s atmosphere at sea level. The photosphere is opaque (not transparent), because it contains negative hydrogen ions (a hydrogen atom with two electrons, instead of the usual one). Hydrogen ions block, absorb, and emit light, all of which prevent light from passing directly through a cloud of hydrogen ions. Some images of the Sun suggest that its white-hot disk is perfectly round and smooth, without a blemish. This uniform appearance is misleading. Under close inspection with a telescope, the photosphere breaks into a million tiny bright points, called granules, with a strongly textured and varying pattern. The hot granules are about 1,500 km (about 900 mi) across, and they are grouped into much larger supergranules about 30,000 km (about 20,000 mi) in diameter. Granules are places within the photosphere where hot, and therefore bright, material reaches the surface. The granules are in constant turmoil and change. Hot gas rises up, liberating its energy. After the gas cools, it sinks downward along the dark lanes between the granules. Each bright cell lasts only a few minutes before it is replaced by another. This honeycomb of rising and falling gas marks the top of the convective zone. Large, dark spots, called sunspots, are often visible in the photosphere. The biggest sunspots exceed Earth in size and are easily visible with a telescope. Sunspots rotate with the Sun and change in size and shape. They come and go, with lifetimes lasting from hours to months. The number of sunspots increases, then decreases, over an 11-year cycle. The position of sunspots changes as the number changes. Sunspots are concentrated in two belts, one north and one south of the solar equator. When the number of sunspots is at a minimum, the belts are near the equator. When the number of sunspots is at its maximum, the belts are at higher latitudes, nearer the poles. Sunspots are places in the Sun’s photosphere that contain magnetic fields thousands of times stronger than Earth’s magnetic field. Sunspots appear dark, because they are much cooler than their bright surroundings. The concentrated magnetism in sunspots keeps them cold. The strong magnetic field of a sunspot acts as a valve, choking off the heat, light, and energy flowing outward from the solar interior. This valvelike action keeps sunspots at a temperature of 3230°C (5850°F), or just over half the temperature of the surrounding photospheric gas. While sunspots are darker than their surroundings, they still radiate light. A sunspot is about ten times brighter than the full Moon. Scientists were perplexed for decades over what holds sunspots together. Scientists believed that the outward pressure of the strong, localized magnetic fields that are concentrated in sunspots should make the sunspots expand and disperse. By examining motions beneath sunspots, helioseismologists have shown that flows of gas converge below sunspots. The converging flows force the surface magnetic fields together and concentrate them to form sunspots. Sunlight appears yellowish, but it is actually a combination of a rainbow of colors. Scientists use special instruments called spectrographs to separate sunlight out into its different colors. These instruments do the same thing that water molecules in the atmosphere do when the molecules produce a rainbow. Each color corresponds to a different wavelength of light. Red has the longest wavelength of visible light, and violet has the shortest. The range of wavelengths of sunlight and the intensity at each wavelength are called the Sun’s spectrum. The study of the spectra of the Sun and other objects or materials is called spectroscopy. When sunlight is spread out like a rainbow in the Sun’s spectrum, many dark gaps separate one color from another in the row of colors. These gaps are called absorption lines. Each absorption line is created when sunlight passes through the gases in the Sun’s photosphere. Atoms and ions of each element in the gas absorb light at certain wavelengths, creating dark gaps in the Sun’s spectrum. The dark absorption lines in the spectra of the Sun and other stars fingerprint the ingredients of these stars. Each chemical element produces a unique set of lines, and the presence of these lines shows that a particular element is present in the stellar photosphere. Darker absorption lines indicate greater absorption and therefore larger amounts of the element. Absorption lines in the Sun’s spectrum show that hydrogen is by far the most abundant element in the Sun. Other prominent absorption lines are produced by helium, sodium, calcium, and iron. Altogether, 92.1 percent of the atoms in the Sun are hydrogen atoms, 7.8 percent are helium atoms, and the other, heavier elements—sodium, calcium, iron, and other elements—make up only 0.1 percent of the atoms in the Sun. The Sun’s absorption lines are called Fraunhofer lines, named after German physicist Joseph von Fraunhofer, who cataloged them in the 1800s. The most common Fraunhofer lines are listed below, by the letter Fraunhofer gave them, the color that they block, and the element that causes them. The Fraunhofer lines designated A and B actually have nothing to do with the composition of the Sun. They only appear on spectra gathered within Earth’s atmosphere. Earth’s atmosphere absorbs sunlight at the wavelengths of the A and B Fraunhofer lines, creating dark lines on the Sun’s spectrum. A spectrum gathered above Earth’s atmosphere would not have these lines. The chromosphere is a thin layer about 2,000 to 3,000 km (about 1,200 to 1,900 mi) thick, just above the visible photosphere. The chromosphere’s temperature rises from 5510°C (9950°F) near the photosphere to about 9700°C (17,500°F) near the corona. At temperatures such as those in the chromosphere, hydrogen emits a distinctive deep red color. Scientists often study the chromosphere by filtering out all sunlight except the light that has the wavelength produced by hydrogen in the chromosphere. Calcium ions (calcium atoms with one electron missing) also produce distinctive radiation in the chromosphere. Calcium ions emit ultraviolet light, or radiation with a wavelength just shorter than visible light. The radiation released by calcium ions is also useful for examining details in the chromosphere. Hydrogen and calcium emissions reveal huge regions of cool, dense gas suspended above the photosphere by powerful magnetic fields. The cool gas looks dark against the brightness of the Sun beneath it. At the edge of the disk of the Sun, where the chromosphere extends beyond the lower layers of the Sun, the gas of the chromosphere creates bright loops called prominences against the dark sky. Against the surface of the Sun, however, the prominences look dark. Prominences are often called filaments when they appear against the background of the hot Sun. Sunspots extend from the photosphere into the chromosphere, creating even darker spots on the chromosphere. Hot gas from the photosphere penetrates the chromosphere around the sunspots, creating bright regions called plages. The corona is the very hot layer of the solar atmosphere above the chromosphere. It extends to Earth and beyond as the solar wind. The Sun’s temperature rises to 2 million degrees C (4 million degrees F) at the bottom of the corona, and remains almost that hot as it reaches Earth. The high temperature of the corona presents one of the most puzzling problems of solar physics. The chromosphere and photosphere are closer to the Sun’s core than is the corona, but the corona is several hundred times hotter than the chromosphere and photosphere. According to the laws of thermodynamics (the branch of physics that deals with the movement and transfer of heat), heat cannot move from a cooler area to a warmer area. Scientists believe that the temperature of the corona results from effects of the Sun’s magnetic fields instead of radiation from the Sun’s core. Comparisons of the corona and the Sun’s magnetic fields have shown that the corona is hottest where the magnetic fields are strongest. The entire corona is stitched together by thin, bright, magnetized loops of material that constrain the hot, dense gas of the corona and shine brightly at X-ray wavelengths. These loops are in a continuous state of change—they can rise from inside the Sun, sink back down into it, or expand into space. They often come together, sometimes merging with each other and sometimes destroying each other. The magnetic loops store magnetic energy. When they interact, the magnetic loops release their stored energy into the corona, providing the energy that keeps the corona so hot. The corona’s magnetic field also has gaps in it, called coronal holes. When astronomers use X-ray telescopes to look at the corona, coronal holes appear as large dark areas, because they are cooler and contain less material than the rest of the corona. Spectral lines come from atoms emitting and absorbing light when their electrons gain or lose energy. The corona is so hot that atoms in the corona are stripped of some of their electrons. These atoms then have different numbers and arrangements of electrons from atoms in the rest of the atmosphere and thus produce different spectral lines. The corona emits most of its radiation at very short ultraviolet and X-ray wavelengths. The underlying photosphere emits very little radiation in these parts of the spectrum, so an image of the Sun in short ultraviolet and X-ray wavelengths produces an accurate picture of the corona. Much of the ultraviolet and X-ray radiation that hits Earth’s atmosphere is absorbed by atoms and molecules in the atmosphere, so scientists use instruments in space to study the corona. Studies of the corona reveal dramatic, violent events called solar flares and coronal mass ejections (CMEs). Solar flares release energy from magnetic loops in the corona, heating the gases of the corona and sending particles and radiation out into the solar system. A coronal mass ejection occurs when an explosion in the corona pushes millions or billions of metric tons of material out into space. The frequency of occurrence of both solar flares and CMEs follows the pattern of the 11-year sunspot cycle (as the number of sunspots increases, so does the number of solar flares and CMEs). Both kinds of solar explosions seem to result from the sudden release of energy stored in coronal magnetic fields. The Sun’s ever-changing magnetism produces unrest on an awesome scale. The sudden, brief, intense outbursts called solar flares can rip through the Sun’s atmosphere with tremendous violence. They release energy equivalent to that of billions of hydrogen bombs in a just few minutes, increasing the temperature of Earth-sized regions of the corona by ten times and flooding the solar system with intense radiation. During a solar flare, the tops of magnetized coronal loops release energy. In less than a second, electrons and positive ions within these loops accelerate to nearly the speed of light. The explosion hurls the electrons and ions out into space and down into the Sun. The particles strike the dense chromosphere below and produce high-energy X rays and gamma rays. Solar flares are probably triggered when oppositely directed magnetic fields come together in the corona, releasing their stored magnetic energy in a manner similar to that of a tightly twisted rubber band that suddenly snaps. After releasing their pent-up energy, the magnetic fields reconnect and relax to a stable configuration. Coronal mass ejections are giant magnetic bubbles that expand to nearly the size of the Sun itself as they leave the low corona. The CMEs move outward at speeds from 200 to 1,000 km/s (100 to 600 mi/s). They carry up to 10 billion metric tons of coronal material into the space of the solar system. They accelerate and propel ahead of them vast quantities of high-speed particles. CMEs sometimes occur when part of the coronal magnetic field becomes sheared and twisted, often disrupting a filament (a loop of material in the chromosphere, also called a prominence). The filament shoots through the chromosphere into the corona, carrying material with it. Earth is affected by the radiation and particles that solar flares and coronal mass ejections release. Intense radiation from a solar flare reaches Earth’s atmosphere in just eight minutes. The X-ray radiation of flares strips electrons from atoms and molecules in Earth’s atmosphere, changing the electrical properties of the atmosphere. This change can disrupt radio communications and make the atmosphere expand farther into space than usual. Friction can develop between the expanded atmosphere and satellites that orbit near Earth, slowing down the satellites. Frequent solar flares can also increase levels of ultraviolet radiation in the atmosphere, which in turn changes oxygen molecules into ozone (oxygen made up of molecules containing three oxygen atoms instead of the usual two). This added ozone actually helps block harmful radiation from the Sun. Particles that solar flares and CMEs release take a day or more to reach Earth. Blasts of these particles can compress Earth’s magnetic field. Disruptions in Earth’s magnetic field can cause geomagnetic storms. Geomagnetic storms occur when Earth’s magnetic field compresses and intensifies, then relaxes back to its normal intensity. The increased intensity of the magnetic field can interfere with signals passing through the atmosphere and cause power surges on wires that carry electricity. CMEs can also trigger intense auroras, colorful displays of light that occur in the atmosphere near Earth’s poles when energetic particles enter the atmosphere. In this case, energetic charged particles collide with atoms and molecules of the atmosphere. This boosts the atoms and molecules to higher energies and forces them to glow. Particles released by a CME can damage or destroy Earth-orbiting satellites and may endanger astronauts in space. Solar flares and CMEs have such a large potential for affecting Earth that space weather forecasters continuously monitor the Sun from ground and space to warn of threatening solar activity. If humans can learn to predict these violent events by pinpointing magnetic changes on the Sun, these predictions will provide very useful early warnings. Flares and CMEs are tied to the cycle of solar activity. The most recent maximum of solar activity occurred in 2001, and the next should occur in 2012. Forecasters study the Sun carefully during these periods. The outermost part of the Sun is a stream of particles that flows from the Sun into the solar system. This part of the Sun, called the solar wind, is the corona expanding into space. The solar wind extends all the way to the heliopause, far past the orbit of Pluto. The corona is so hot that it cannot stand still. It is expanding outward in all directions, filling the solar system with a ceaseless flow of electrons, ions, and magnetic fields. The solar wind has two components. The fast part of the wind pours out of the regions near the poles of the Sun at speeds around 750 km/s (around 470 mi/s). The slower component of the solar wind gusts unevenly from the Sun’s equatorial regions at speeds from 300 to 400 km/s (190 to 250 mi/s). Scientists believe that the fastest part of the solar wind leaves the Sun through coronal holes, cool spots in the corona. The magnetic field of the Sun is relatively weak around coronal holes and thus allows particles in the solar wind to escape. Heavier particles seem to move more quickly than lighter particles in the same stream within coronal holes. The intermittent gusts from nearer the equator come from solar flares and coronal mass ejections. Both components of the solar wind gain speed as they spread out and leave the Sun. The fast component reaches its top speed close to the Sun, but the slow solar wind continues gaining speed much farther out. The Sun rotates as it emits the solar wind, so the solar wind spirals around the solar system. The solar wind carries the Sun’s magnetic field with it and sets up a spiral magnetic field throughout the solar system. The solar wind and its magnetic field affect the magnetic fields of the planets, the direction of the tails of comets, and even the flight paths of spacecraft. The Sun is so important to life on Earth that humans have always paid special attention to it. The movement of the Sun across the sky helps mark time. The change throughout the year in the Sun’s daily path helps mark the seasons. Many cultures attach special significance to solar events, such as eclipses. The brightness of the Sun made studying it closely difficult for humans for many years. Looking at the Sun directly is dangerous, and even thick clouds do little to protect human eyes from the damage that direct sunlight causes. Astronomers could not make true scientific studies of the Sun until they developed techniques to observe the Sun indirectly. The study of the Sun has both pushed and been pushed by revolutionary scientific discoveries. Early indirect observations of the Sun, using a telescope, allowed scientific study of the Sun to begin, showing that the Sun is a dynamic, changing body. The development of spectroscopy and the discovery of elementary particles and nuclear fusion allowed scientists to begin to understand the composition of the Sun and the processes that fuel it. The development of artificial satellites and other spacecraft finally allowed scientists to study the Sun from space, allowing a full view of all of the Sun’s radiation and a continuous study of the Sun. Greek philosopher Aristotle was the first known person to use a device that allowed indirect observation of the Sun. Sometime between 384 bc and 322 bc Aristotle noticed that a hole in a screen would create an image of the Sun on the ground, if the screen were between the Sun and the ground. He made a simple version of a device called a camera obscura to take advantage of this effect. A camera obscura is still a popular way to observe solar eclipses. Italian scientist Galileo observed the Sun with a telescope for the first time in 1610. Looking through a telescope directly at the Sun is even more dangerous than looking at the Sun with the naked eye, so Galileo turned the telescope into a camera obscura. He pointed it at the Sun and then set up a screen behind the eyepiece. The eyepiece projected the image of the Sun onto the screen. Galileo observed sunspots with his telescope. He saw that sunspots rotate with the Sun and change in size and shape. Galileo’s work showed that the Sun is a changing and active body. The next major breakthrough in the study of the Sun was the development of ways to study sunlight. In the mid-17th century English scientist Isaac Newton used a prism—a specially cut chunk of glass—to break sunlight down into its different colors. This range of colors is called the Sun’s spectrum, and the study of spectra is called spectroscopy. In 1802 British scientist William Wollaston found that the solar spectrum was cut by several dark gaps. By 1815 German physicist Joseph von Fraunhofer had cataloged the wavelengths of more than 300 of the gaps, called absorption lines. Fraunhofer assigned letters to the most prominent absorption lines. In the mid-19th century German scientists Gustav Kirchhoff and Robert Bunsen related the absorption lines in the Sun’s spectrum to chemical elements. In 1925 English-born American astronomer Cecilia Payne (later Cecilia Payne-Gaposchkin) compared the spectrum of the Sun to that of other stars to show that virtually all bright, middle-aged stars have the same composition. The spectrum of the Sun’s corona was studied for the first time in the mid-19th century. During the solar eclipse of August 7, 1869, American astronomers Charles A. Young and William Harkness independently discovered that the corona’s spectrum featured an especially bright line of green light. Bright lines in a spectrum are called emission lines. They are the fingerprints of elements in the substance producing the light. The corona’s bright green emission line comes from highly ionized iron, indicating that the corona has very high temperatures. Detailed studies of the Sun’s photosphere and the sunspots began with Galileo’s telescopic camera obscura of the 17th century. The next revolution in this area occurred in the 1840s, when German scientist Heinrich Schwabe discovered that the number and positions of sunspots vary over an 11-year period. In 1859 British astronomer Richard Carrington discovered solar flares. Carrington’s discovery helped explain that geomagnetic storms (increased intensity of Earth’s magnetic field) are related to events on the Sun. In 1908 American astronomer George Ellery Hale showed that sunspots contain magnetic fields that are thousands of times stronger than Earth’s magnetic field. The Sun produces an enormous amount of energy. Scientists could not explain how something with the mass of the Sun could produce so much energy until they discovered nuclear fusion. The details of just how nuclear fusion changes hydrogen into helium nuclei were not known until discoveries in the field of elementary particles were made. Elementary particles are the tiny particles that make up all matter. The most familiar particles, the particles that make up atoms, are protons, neutrons, and electrons. Protons and neutrons are the main particles involved in nuclear fusion. Both types of particles are about the same size and mass, but protons have a positive electric charge, while neutrons are electrically neutral. New Zealand-born British physicist Ernest Rutherford discovered the proton in 1918. British physicist Sir James Chadwick discovered the neutron in 1932, and was awarded the 1935 Nobel Prize in physics for his discovery. The first fusion reaction in a laboratory occurred in the early 1930s. In 1938 German-born American physicist Hans A. Bethe and American physicist Charles L. Critchfield demonstrated how a sequence of nuclear reactions, called the proton-proton chain, makes the Sun shine. Bethe was awarded the 1967 Nobel Prize in physics for his discoveries concerning energy production in stars. After scientists understood how the Sun produces its energy, they began developing theories to explain how the Sun’s energy travels from the core to the Sun’s atmosphere. For the first few decades after the discovery that fusion powers the Sun, scientists deduced the Sun’s structure by comparing the theoretical output of the Sun’s core to the energy actually released at the Sun’s atmosphere. In the 1960s American physicist Robert Leighton developed a camera that could record Doppler shifts in light at the Sun’s surface. A Doppler shift is a change in the wavelength of light caused by the movement of the object that is emitting the light. If the object is moving away from the observer, each wave will have to travel farther to reach the observer, making the distance between waves (the wavelength) longer. An object moving toward the observer will seem to emit light with a shorter wavelength. Leighton used this device to discover that the Sun seemed to pulsate in and out, making a complete cycle about every five minutes. See the Oscillating Sun section of this article. Leighton’s discovery launched the field of helioseismology, or the study of the Sun’s interior by observing the vibrations of the Sun and how sound waves move through it. In the 1970s scientists demonstrated that the entire Sun is vibrating with ponderous, organized rhythms that can extend to its very core. Scientists developed models of the interior of the Sun based on vibrations at its surface. In 1995 six observatories around the world coordinated with each other to begin observing the oscillations of the Sun as a team. This project, a collaboration of 20 nations, is called the Global Oscillation Network Group (GONG). GONG can keep constant watch on the Sun because, at any given time, daytime is being experienced by at least one of the observatories. GONG has allowed scientists to get a better idea of the interior structure of the Sun through helioseismology. Studying the Sun from space has revolutionized solar physics. Since the first observations from space began to be made, scientists have made great advances in the study of the Sun’s surface, energy production, structure, and relationship to Earth and the solar system. Study from space began in 1957 when the Soviet satellite Sputnik 2, the second satellite to go into space, carried instruments to study the Sun. Since then, many missions have been devoted to studying the Sun. The series of missions by the Pioneer spacecraft of the United States included several experiments to study the Sun and its relationship to Earth. The Pioneer program lasted from the late 1950s to the 1970s. The U.S. Mariner 2 spacecraft, launched in 1962, used data obtained from its voyage to Venus to demonstrate that a low-speed solar wind is continuously emitted from the Sun, and also discovered high-speed streams in the Sun’s winds. In the 1960s and 1970s the U.S. Orbiting Solar Observatory (OSO) series studied the Sun over an entire cycle of solar activity. One of the satellites in the OSO series was the piloted Skylab space station, launched in 1973. Skylab astronauts used X-ray telescopes to transform our knowledge of the Sun’s corona. In the early 1980s the United States launched the Solar Maximum Mission spacecraft to study the Sun during its most active period. The joint Japanese, U.S., and British probe Yohkoh studied solar flares through the 1990s. Two of the most productive solar spacecraft of the late 1990s and early 2000s were the Ulysses spacecraft and the Solar and Heliospheric Observatory (SOHO). Both spacecraft are joint projects of the United States and the European Space Agency (ESA). Ulysses was launched in 1990 and SOHO was launched in 1995. Ulysses’s orbit takes it over the poles of the Sun, then back to the planet Jupiter to get a gravitational boost that sends the spacecraft back to the Sun. By 2001 Ulysses had passed over the Sun’s north and south poles twice. Ulysses’s mission has contributed much knowledge about the solar wind in regions above the Sun’s poles. Findings from this mission conclusively demonstrated that the fast component of the solar wind pours out at high solar latitudes, including from polar coronal holes. The slow component of the solar wind is constrained to low latitudes near the solar equator. Ulysses also found that the Sun’s magnetic field does not warp at the poles as much as scientists expected. SOHO is at a point in space where the Sun’s gravitational pull balances Earth’s gravitational pull, so the satellite orbits the Sun with Earth. SOHO always faces the Sun. This probe has allowed scientists to make great leaps forward in their understanding of the structure and dynamics of the solar interior, the heating mechanisms of the solar corona, and the origin and acceleration of the solar wind. SOHO has returned amazing images of the Sun, including comets hitting the Sun and features on the Sun’s surface that scientists compare to tidal waves, tornadoes, and rivers.