Solar structure and energy

Advances in both theoretical physics and astronomical observing capabilities have provided an increasingly detailed description of the Sun. As the nearest star to us, the Sun can also function as a laboratory for investigations of stellar physics in general.

Overview

The Sun is composed mostly of hydrogen. When the Sun first formed, there was enough hydrogen in the estimated size of its core for this reaction to keep the Sun shining for approximately ten billion years. Since the Sun is about 4.5 billion years old, it is only about halfway through generating energy in its core by hydrogen fusion.

The Sun’s mass is about 2 10³⁰ kilograms. It consists primarily of the two simplest atoms: hydrogen (about 91 percent by number of atoms, about 71 percent by mass) and helium (about 9 percent by number of atoms, about 27 percent by mass). Other chemical elements account for a combined total of about 0.1 percent by number of atoms or 2 percent by mass. More than sixty chemical elements have been identified in the Sun’s spectrum, and probably all the elements are present in minute amounts. Some of the more abundant of the other elements include oxygen, carbon, nitrogen, silicon, magnesium, neon, iron, and sulfur.

A major clue to the conditions inside the Sun is that it seems to be reasonably stable, with no large, rapid changes. Thus, it must be in a state of hydrostatic equilibrium. This means that its self-gravity, which tries to make it contract, is balanced by its internal pressure, which tries to make it expand. If it were not in hydrostatic equilibrium (or at least very, very close to it), the Sun would either be contracting or expanding noticeably.

The Sun’s luminosity—that is, the rate at which it emits electromagnetic energy—is approximately 3.8 10²⁶ joules per second. The nuclear reaction that supplies the Sun’s energy needs is the fusion of four hydrogen nuclei into one helium nucleus. In the production of this single helium nucleus, 4.8 10^-29 kilograms of mass are converted into 4.3 10^-12 joules of energy. Thus, to generate the Sun’s luminosity (3.8 10²⁶ joules per second), every second, about 36 10³⁷ hydrogen nuclei (with a total mass of 602,300,000 metric tons) are fused into 9 10³⁷ helium nuclei (with a total mass of 598,100,000 metric tons), and the excess 4,200,000 metric tons of mass is converted into energy in the Sun’s core.

From the center outward, the Sun has several layers or zones: the hydrogen fusion core, the radiative zone, the convective zone, the photosphere, the chromosphere, and the corona. However, other than the photosphere, these layers are either difficult or impossible to observe. Above the photosphere, the gas is so tenuous that it emits much less visible light. Below the photosphere, the gas is opaque—sort of like fog—so electromagnetic radiation cannot escape. Consequently, the layers below the photosphere must be studied indirectly via computer models of the Sun’s interior created from images. The layers above the photosphere can be seen visually on special occasions or observed anytime at various wavelengths outside the visible range.

The energy generated in the Sun’s core initially makes its way outward by radiation—that is, as a flow of photons. However, because the Sun’s interior is so opaque, a photon travels only about one centimeter before it is absorbed and reemitted by an atom. This occurs over and over again, so this method of energy transport is called radiative diffusion. The region of the Sun in which energy is transported entirely by radiative diffusion (or simply radiation) is its radiative (or radiation) zone, and this zone extends from the center out to about two-thirds of the Sun’s radius.

Beyond this distance, the gas becomes even more opaque, making it more difficult for photons to travel. As a result, convection sets in, and energy is carried by rising bubbles of hot, lower-density gas. After transferring their heat, the cooler, denser bubbles sink. The region of the Sun in which convection occurs is its convective (or convection) zone, and it occupies the outer third of the Sun’s radius. Even in this zone in which convection occurs, some energy is also carried by radiation, that is, by photon flow.

The next layer, the photosphere or visible surface of the Sun, is the lowest level of the solar atmosphere, with a radius of about 700,000 kilometers and a thickness of about 500 kilometers. Temperatures decrease from about 6,600 kelvins at its base to about 4,400 kelvins at its top; the overall effective temperature of the photosphere is 5,800 kelvins.

When examined with high-resolution imagers, the photosphere reveals a wealth of structure and detail. Most pronounced is the presence of granulation, an alternation of brighter spots with darker borders resembling a mixture of salt and pepper. Each granule is a region of gas about 1,000 kilometers in diameter, larger than the state of Texas. The photosphere is at the top of the Sun’s convective zone, the part of the Sun’s interior where energy is transported outward by convection. The brighter region at the center of the granule is a hot gas bubble rising, while the darker border regions of a granule are cooler gases sinking back down. The convection zone in the Sun produces waves of thermal energy that shoot up through the photosphere. These waves make the photosphere appear to oscillate, with periods ranging from minutes to hours.

Above the photosphere are two more layers of the solar atmosphere, called the chromosphere and the corona, that usually cannot be seen with the unaided eye. The photosphere emits so much light and is so much brighter than the rest of the solar atmosphere that it usually overwhelms the weaker visible light from the layers above it, making them difficult to see. During a total solar eclipse, however, the Moon blocks the bright photosphere, and the fainter chromosphere and corona are visible to the unaided eye, extending out around the silhouette of the Moon.

Chromosphere literally means color sphere.”When seen during total solar eclipses, it appears as a narrow red ring just beyond the Moon’s silhouette. The red color is one of the wavelengths (656.3 nanometers) that can be strongly emitted and absorbed by hydrogen atoms. The thickness of the chromosphere is about 2,000 kilometers, with temperatures of 4,500 kelvins at its base and rising to 30,000 kelvins at its top. These temperatures result in strong ultraviolet emission, and the chromosphere can be observed at any time (not just during eclipses) at ultraviolet wavelengths. However, our atmosphere is opaque to most of the ultraviolet, so ultraviolet observations must be conducted by spacecraft above our atmosphere.

The chromosphere contains hundreds of thousands of thin spikes called spicules. These spicules are jets of hot gas, hundreds of kilometers across and thousands of kilometers tall. They rise dramatically and then fall over a lifetime of several minutes, thus creating a dynamic appearance, like the dance of many small candle flames. The chromosphere also has a granulated structure. This structure cannot be directly observed but has been deduced from spectroscopic studies of the motions of gas in the chromosphere using the Doppler effect (a change in the frequency and wavelength of electromagnetic radiation caused by the motion of its source toward or away from the observer). Such studies have revealed that the chromosphere contains large, organized cells of gas called supergranules that move in unison under the influence of convective forces. These supergranules are 30,000 kilometers in diameter and contain hundreds of normal granulation regions.

Above the chromosphere, the density rapidly decreases, and the temperature abruptly increases to about one million kelvins in a thin layer called the transition region. Beyond is the corona, the outermost part of the solar atmosphere. The corona’s density is so low that it emits relatively little visible light, but during a total solar eclipse, the corona can be seen as a broad, glowing white halo of light out to distances of several solar radii. The temperature of the corona is about one million to two million kelvins. Consequently, it strongly emits X-rays. It can be observed at any time at X-ray wavelengths, but our atmosphere is opaque to X-rays (just as it is opaque to most ultraviolet radiation), so X-ray observations of the corona must also be conducted from spacecraft above our atmosphere.

Applications

Our Sun is a star, the only star in our solar system. It is often described as an average star. While it is average in the sense that it is a “main sequence star” and just about in the middle of the ranges of stellar luminosity, mass, diameter, surface temperature, and various other properties, it is not typical in the sense that the vast majority of stars are smaller, cooler, and less luminous than the Sun. Furthermore, the comparatively small number of stars that are larger and more luminous than the Sun are so bright that they account for most of the light emitted within a galaxy of billions of stars.

Nevertheless, the Sun is the closest star to us—at a distance of about 150 million kilometers (which defines one astronomical unit, or one AU)—and, as such, provides a “laboratory” where normal stellar processes can be observed and studied at relatively close range. The next closest star, Proxima Centauri (which probably is a member of the Alpha Centauri star system and thus sometimes is called Alpha Centauri C), is about 40 trillion kilometers away, or about 270,000 times the distance to the Sun. Light from the Sun takes only five hundred seconds (a little over eight minutes) to reach Earth, while light from Proxima Centauri takes more than four years to reach us.

Aside from the fact that most other stars are smaller and less luminous, they do share have much in common with the Sun. Most stars have compositions similar to the Sun’s—mostly hydrogen, some helium, and very small amounts of other chemical elements. Most stars, like the Sun, are reasonably stable and thus are in hydrostatic equilibrium. The light we receive from other stars, like the light we receive from the Sun, comes from a layer that resembles the photosphere (and hence, other stars are said to possess “photospheres”). Observational evidence suggests that at least some stars (and probably most, if not all) possess chromospheres and coronas as well—like the Sun.

Also like the Sun, stars generate energy for most of their active “lives” by nuclear fusion reactions. The fusion of hydrogen into helium in its core is the first and by far the longest-lasting fusion reaction that a star can tap. Most stars thus, are generating energy the same way the Sun is, by hydrogen-to-helium fusion in their cores. After a stellar core’s hydrogen is exhausted, other nuclear fusion reactions are possible, such as the fusion of three helium nuclei into a single carbon nucleus, but these higher fusion reactions do not supply energy for long.

The energy generated in a star’s core is transported outward in most stars by the same two processes operating in the Sun: radiation and convection. However, computer models of stellar interiors indicate that in some stars, the convective zone is in the central part of the star, and the radiative zone is farther out. Furthermore, in some parts of some stars in some stages of their lives, a third energy-transport mechanism, conduction, is more effective than either radiation or convection.

Context

Unraveling the mystery of how the Sun and most other stars shine is one of the great achievements of science in the twentieth century. Since Anaxagora's assertion in the fifth century BCE that the Sun was not a god but simply a big ball of fire, the question of solar fuel was a mystery. An early suggestion for the energy source was some form of chemical combustion; in other words, the Sun was literally a ball of fire that was burning something, just as Anaxagoras said. The problem was the timescale. No matter what the Sun was made of, it would have to chemically consume (burn) its entire mass in a time on the order of thousands to tens of thousands of years to keep shining at its present rate.

Another suggestion was meteoritic infall; that is, the Sun was heated by matter falling into it. Again, there is a problem of scale. A rocky asteroid eleven kilometers in diameter would have to slam into the Sun every second to release enough energy to supply its present luminosity. Also, such a rate would have produced some observable immediate effects, but none had ever been detected.

By the middle of the nineteenth century, the answer seemed to be slow gravitational contraction. The Sun would only have to shrink by a few tens of meters each year (an insignificant amount compared to its radius of 700,000 kilometers) to provide for its present luminosity. Furthermore, slow gravitational contraction could keep the Sun shining for tens of millions of years. At the time, it seemed like a reasonable duration to astronomers and physicists, but geologists, paleontologists, and biologists were beginning to figure out that the Earth and its rock and fossil record were at least hundreds of millions, if not billions of years old. Again, the timescale did not fit the theory.

The accepted energy source now is nuclear fusion, a process in which light atomic nuclei are fused into heavier atomic nuclei, releasing energy. The mass of the resulting nucleus is not quite as great as the sum of the masses of the nuclei that fused to form it. That small difference in mass gets converted into energy according to Einstein’s famous equation, E = mc², which says that matter, m, and energy, E, are equivalent and are related by a physical constant (the speed of light, c, squared). To be able to overcome the electrostatic repulsion that positively charged nuclei feel toward each other, nuclear fusion reactions can occur only at high temperatures (so the nuclei are moving rapidly) and high densities (so the nuclei are close together), exactly the conditions at the center of the Sun and most stars. According to computer models of the Sun’s interior, the central temperature is about fifteen million kelvins and the central density is about 150 times the density of water.

Our understanding of the structure and physics of the Sun has been obtained in two ways: through observations of the Sun in all parts of the electromagnetic spectrum and through computer models that calculate conditions inside the Sun. The computer models are put together through analyzing images taken of the Sun by devices such as the Solar Orbiter, a joint effort between the European Space Agency and NASA. The processes that take place in the Sun are representative of processes that take place in stars generally. Thus, our solar research has also contributed to our study of stars throughout our Milky Way galaxy and in galaxies beyond the Milky Way.

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