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Chapter 12

Chapter 12. How the Stars Shine: Cosmic Furnaces. Introduction. E ven though individual stars shine for a relatively long time, they are not eternal. Stars are born out of the gas and dust that exist within a galaxy; they then begin to shine brightly on their own. Eventually, they die.

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Chapter 12

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  1. Chapter 12 How the Stars Shine: Cosmic Furnaces

  2. Introduction • Even though individual stars shine for a relatively long time, they are not eternal. • Stars are born out of the gas and dust that exist within a galaxy; they then begin to shine brightly on their own. • Eventually, they die. • Though we can directly observe only the outer layers of stars, we can deduce that the temperatures at their centers must be millions of kelvins. • We can even figure out what it is deep down inside that makes the stars shine. • To determine the probable life history of a typical star, we observe stars having many different ages and assume that they evolve in a similar manner. • However, we must take into account the different masses of stars; some aspects of their evolution depend critically on mass.

  3. Introduction • We start this chapter by discussing the birth of stars. • We see how new capabilities of observing in the infrared in addition to the visible are helping us understand star formation (see figure). • We then consider the processes that go on inside a star during its life on the main sequence. • Finally, we begin the story of the evolution of stars when they finish the main-sequence stage of their lives. • Chapters 13 and 14 will continue the story of what is called “stellar evolution,” all the way to the deaths of stars.

  4. Introduction • Near the end of this chapter we will see that the most important experiment to test whether we understand how stars shine is the search for elusive particles, called neutrinos, from the Sun. • Over the past decades a search for them has been made, but only about a third to half of those expected had been found. • Recent experiments have provided better ways of detecting neutrinos than we previously had, and they were there all along, though transformed and thus hidden! • The results indicate that we did not understand neutrinos as well as we had thought. • These astronomical results therefore have added important knowledge about fundamental physics in addition to our understanding of the stars.

  5. 12.1 Starbirth • The birth of a star begins with a nebula—a large region of gas and dust (see figures). • The dust (tiny solid particles) may have escaped from the outer atmospheres of giant stars.

  6. 12.1 Starbirth • The regions of gas and dust (often called clouds, or “giant molecular clouds”) from which stars are forming are best observed in the infrared and radio regions of the spectrum, because most other forms of radiation (such as optical and ultraviolet) cannot penetrate them. • We discuss the infrared observations largely here, including the new capabilities of NASA’s Spitzer Space Telescope, and we leave the radio observations to Chapter 15 on the Milky Way Galaxy. • Stars forming at the present time incorporate this previously cycled gas and dust, which gives them their relatively high abundances of elements heavier than helium in the periodic table (still totaling less than a per cent). • In contrast, the oldest stars we see formed long ago when only primordial hydrogen and helium were present, and therefore have lower abundances of the heavier elements, as we discussed near the end of the previous chapter.

  7. 12.1a Collapse of a Cloud • Consider a region that reaches a higher density than its surroundings, perhaps from a random fluctuation in density or—in a leading theory of why galaxies have spiral arms—because a wave of compression passes by. • Still another possibility is that a nearby star explodes (a “supernova”; see Chapter 13), sending out a shock wave that compresses the gas and dust. • In any case, once the cloud gains a higher-than-average density, the gas and dust continue to collapse due to gravity. • Energy is released, and the material accelerates inward. • Magnetic fields may resist the infalling gas, slowing the infall, though the role of magnetic fields is not well understood in detail.

  8. 12.1a Collapse of a Cloud • Eventually, dense cores, each with a mass comparable to that of a star, form and grow like tiny seeds within the vast cloud. • These protostars (from the prefix of Greek origin meaning “primitive”), which will collectively form a star cluster, continue to collapse, almost unopposed by internal pressure. • But when a protostar becomes sufficiently dense, frequent collisions occur among its particles; hence, part of the gravitational energy released during subsequent collapse goes into heating the gas, increasing its internal pressure. (In general, compression heats a gas; for example, a bicycle tire feels warm after air is vigorously pumped into it.)

  9. 12.1a Collapse of a Cloud • The rising internal pressure, which is highest in the protostar’s center and decreases outward, slows down the collapse until it becomes very gradual and more accurately described as contraction. • The object is now called a pre-main-sequence star (see figure).

  10. 12.1a Collapse of a Cloud • By this time, the object has contracted by a huge fraction, a factor of 10 million, from about 10 trillion km across to about a million km across—that is, something initially larger than the whole Solar System collapses until most of its mass is in the form of a single star. • During the contraction phase, a disk tends to form because the original nebula was rotating slightly. • We discussed this process when considering the “nebular hypothesis” for the formation of our Solar System (Chapter 9). • Dusty disks have been found around young stars and pre-main-sequence stars in nebulae such as the Orion Nebula (see figure). • These are sometimes called protoplanetary disks (“proplyds”), and they support the theoretical expectation that planetary systems are common.

  11. 12.1a Collapse of a Cloud • Jets of gas are commonly ejected in opposite directions out the poles of the rotating pre-mainsequence star (see figure, right). • As energy is radiated from the surface of the pre-main-sequence star, its internal pressure decreases, and it gradually contracts. • This release of gravitational energy heats the interior, thereby increasing the internal temperature and pressure. • It is also the source of the radiated energy. • Gravitational energy was released in this way in the early Solar System. • As the temperature in the interior rises, the outward force resulting from the outwardly decreasing pressure increases, and eventually it balances the inward force of gravity, a condition known as “hydrostatic equilibrium.” • As we shall discuss later, this mechanical balance is the key to understanding stable stars; see Figure at left.

  12. 12.1a Collapse of a Cloud • Theoretical analysis shows that the dust surrounding the stellar embryo we call a pre-main-sequence star should absorb much of the radiation that the object emits. • The radiation from the pre-main-sequence star should heat the dust to temperatures that cause it to radiate primarily in the infrared. • Infrared astronomers have found many objects that are especially bright in the infrared but that have no known optical counterparts. • These objects seem to be located in regions where the presence of a lot of dust, gas, and young stars indicates that star formation might still be going on.

  13. 12.1a Collapse of a Cloud • Imaging in the visible (with the Hubble Space Telescope) and in the infrared (not only with previous infrared space telescopes but now especially with the Spitzer Space Telescope) has shown how young stars are born inside giant pillars of gas and dust inside certain nebulae. • The Eagle Nebula is the most famous example because of the beautiful Hubble image showing exquisite detail, with false colors assigned to different filters (see figure).

  14. 12.1a Collapse of a Cloud • As hot stars form, their intense radiation evaporates the gas and dust around them, freeing them from the cocoons of gas and dust in which they were born. • We see this “evaporation” taking place at the tops of the Eagle Nebula’s “pillars.” • The stars are destroying their birthplaces as they become independent and more visible from afar.

  15. 12.1b The Birth Cries of Stars • To their surprise, astronomers have discovered that young stars send matter out in oppositely directed beams, while they had expected to find only evidence of infall. • This “bipolar ejection” (see figure) may imply that a disk of matter orbits such premain-sequence stars, blocking an outward flow of gas in the equatorial direction and later coalescing into planets. • Thus the flow of gas is channeled toward the poles.

  16. 12.1b The Birth Cries of Stars • Sometimes clumps of gas appear, but only recently have they been identified with ejections from stars in the process of collapsing. • The clumps seem like spinning bullets, though what makes them spin is uncertain (perhaps connected with the magnetic field). • The ejection of these spinning clumps helps slow the star’s rate of spin, since they carry away angular momentum from what had been a rapidly spinning pre-mainsequence star. • At the same time, some gas with low angular momentum is falling in toward the star. • Hidden here is perhaps the main unsolved problem in star formation at the moment: How do stars figure out what their final mass will be?

  17. 12.1b The Birth Cries of Stars • The bipolar ejection appears as “Herbig-Haro objects,” clouds of interstellar gas heated by shock waves from jets of high-speed gas. • The jets are being ejected from a premain-sequence star, a star in the process of being born. • Since the pre-main-sequence star is hidden in visible light by a dusty cocoon of gas, infrared observations of Herbig- Haro objects most clearly reveal what is going on (see figure).

  18. 12.1b The Birth Cries of Stars • The jets of gas were formed as the pre-main-sequence star contracted under the force of its own gravity. • Because a thick disk of cool gas and dust surrounds the premain-sequence star, the gas squirts outward along the pre-main-sequence star’s axis of rotation at speeds of perhaps 1 million km /hr. • HH–1 and HH–2 (see figure) are more irregular in shape than many other Herbig-Haro objects, perhaps because the bow shock wave we are seeing (a shock wave like those formed by the bow of a boat plowing through the water) has broken up. • These objects are about 1500 light-years from us, in a star-forming region of the constellation Orion. • The smallest features resolved are about the size of our Solar System, and the whole image is only about 1 light-year across.

  19. 12.1b The Birth Cries of Stars • Several classes of stars that vary erratically in brightness have been found. • One of these classes, called T Tauri, contains pre-main-sequence stars as massive as or less massive than the Sun. • Presumably, these stars are so young that they have not quite settled down to a steady and reliable existence on the main sequence. (T Tauri stars always have the word “stars” in their name though technically they haven’t reached the main sequence, so they are not yet fully formed stars.) • In astronomical teaching, we have the question of whether to first consider the formation of stars in the star section of the book, as here, or in the section about the gas and dust between the stars from which the stars form. • We choose to do some of each, and will continue our discussion of stars in formation in that latter location, Chapter 15 on the Milky Way Galaxy.

  20. 12.2 Where Stars Get Their Energy • If the Sun got all of its energy from gravitational contraction, it could have shined for only about 30 million years, not very long on an astronomical timescale. • Yet we know that rocks about 4 billion years old have been found on Earth, and up to 4.4 billion years old on the Moon, so the Sun and the Solar System have been around at least that long. • Moreover, fossil records of planets and animals, which presumably used the Sun’s light and heat, date back billions of years. • Some other source of energy must hold the Sun and other stars up against their own gravitational pull.

  21. 12.2 Where Stars Get Their Energy • Actually, a pre-main-sequence star will heat up until its central portions become hot enough (at least one million kelvins) for nuclear fusion to take place, at which time it reaches the main sequence of the temperature-luminosity (temperature-magnitude, or Hertzsprung-Russell; see Chapter 11) diagram. • Using this process, which we will soon discuss in detail, the star can generate enough energy to support it during its entire lifetime on the main sequence. • A star’s luminosity and temperature change little while it is on the main sequence; nuclear reactions provide the stability.

  22. 12.2 Where Stars Get Their Energy • The energy makes the particles in the star move around rapidly. • Such rapid, random motions in a gas are the definition of high temperature. • The thermal pressure, the force from these moving particles pushing on each area of gas, is also high. • The varying pressure, which decreases outward from the center, produces a force that pushes outward on any given pocket of gas. • This outward force balances gravity’s inward pull on the pocket (“hydrostatic equilibrium,” which we illustrated in the figure).

  23. 12.2 Where Stars Get Their Energy • The basic fusion process in main-sequence stars fuses four hydrogen nuclei into one helium nucleus. • In the process, tremendous amounts of energy are released. (Hydrogen bombs on Earth fuse hydrogen nuclei into helium, but use different fusion sequences. The fusion sequences that occur in stars are far too slow for bombs.) • A hydrogen nucleus is but a single proton. • A helium nucleus is more complex; it consists of two protons and two neutrons (see figure). • The mass of the helium nucleus that is the final product of the fusion process is slightly less than the sum of the masses of the four hydrogen nuclei (protons) that went into it. • A small amount of the mass, m, “disappears” in the process: 0.007 (0.7 per cent) of the mass of the four protons.

  24. 12.2 Where Stars Get Their Energy • The mass difference does not really disappear, but rather is converted into energy, E, according to Albert Einstein’s famous formula E =mc2, where c is the speed of light. • Even though m is only a small fraction of the original mass, the amount of energy released is prodigious; in the formula, c is a very large number. • This energy is known as the “binding energy” of the nucleus, here specifically that of helium. • The loss of only 0.7 per cent of the central part of the Sun, for example, is enough to allow the Sun to radiate as much as it does at its present rate for a period of about ten billion (1010) years. • This fact, not realized until 1920 and worked out in more detail in the 1930s, solved the longstanding problem of where the Sun and the other stars get their energy.

  25. 12.2 Where Stars Get Their Energy • All the main-sequence stars are approximately 90 per cent hydrogen (that is, 90 per cent of the atoms are hydrogen), so there is a lot of raw material to fuel the nuclear “fires.” • We speak colloquially of “nuclear burning,” although, of course, the processes are quite different from the chemical processes that are involved in the “burning” of logs or of autumn leaves. • In order to be able to discuss these processes, we must first review the general structure of nuclei and atoms.

  26. 12.3 Atoms and Nuclei • As we mentioned in Chapter 2, an atom consists of a small nucleus surrounded by electrons. • Most of the mass of the atom is in the nucleus, which takes up a very small volume in the center of the atom. • The effective size of the atom, the chemical interactions of atoms to form molecules, and the nature of spectra are all determined by the electrons.

  27. 12.3a Subatomic Particles • The nuclear particles with which we need to be most familiar are the proton and neutron. • Both of these particles have nearly the same mass, 1836 times greater than the mass of an electron, though still tiny. • The neutron has no electric charge and the proton has one unit of positive electric charge. • The electrons, which surround the nucleus, have one unit each of negative electric charge. • When an atom loses an electron, it has a net positive charge of 1 unit for each electron lost. • The atom is now a form of ion (see figure).

  28. 12.3a Subatomic Particles • Since the number of protons in the nucleus determines the charge of the nucleus, it also dictates the quota of electrons that the neutral state of the atom must have. • To be neutral, after all, there must be equal numbers of positive and negative charges. • Each element (sometimes called “chemical element”) is defined by the specific number of protons in its nucleus. • The element with one proton is hydrogen, that with two protons is helium, that with three protons is lithium, and so on.

  29. 12.3b Isotopes • Though a given element always has the same number of protons in a nucleus, it can have several different numbers of neutrons. (The number of neutrons is usually somewhere between 1 and 2 times the number of protons. The most common form of hydrogen, just a single proton, is the main exception to this rule.) • The possible forms of the same element having different numbers of neutrons are called isotopes. • For example, the nucleus of ordinary hydrogen contains one proton and no neutrons. • An isotope of hydrogen (see figure) called deuterium (and sometimes “heavy hydrogen”) has one proton and one neutron. • Another isotope of hydrogen called tritium has one proton and two neutrons.

  30. 12.3b Isotopes • Most isotopes do not have specific names, and we keep track of the numbers of protons and neutrons with a system of superscripts and subscripts. • The subscript before the symbol denoting the element is the number of protons (called the atomic number), and a superscript after the symbol is the total number of protons and neutrons together (called the mass number, or atomic mass). • For example, 1H2 is deuterium, since deuterium has one proton, which gives the subscript, and an atomic mass of 2, which gives the superscript. (Note that 21H is also correct notation.) • Deuterium has atomic number equal to 1 and mass number equal to 2. • Similarly, 92U238 is an isotope of uranium with 92 protons (atomic number 92) and mass number of 238, which is divided into 92 protons and 23892=146 neutrons. • Each element has only certain isotopes. • For example, most of the naturally occurring helium is in the form 2He4, with a much lesser amount as 2He3.

  31. 12.3c Radioactivity and Neutrinos • Sometimes an isotope is not stable, in that after a time it will spontaneously change into another isotope or element; we say that such an isotope is radioactive. • The most massive elements, those past uranium, are all radioactive, and have average lifetimes that are very short. • It has been theoretically predicted that around element 114, elements should begin being somewhat more stable again. • The handful of atoms of element 114 and 116 discovered in 1998 and 1999 are more stable than those of slightly lower mass numbers— lasting even about 5 seconds instead of a small fraction of a second. (A claim that element 118 was also discovered has been withdrawn.)

  32. 12.3c Radioactivity and Neutrinos • During certain types of radioactive decay, as well as when a free proton and electron combine to form a neutron, a particle called a neutrino is given off. • A neutrino is a neutral particle (its name comes from the Italian for “little neutral one”). • Neutrinos have a very useful property for the purpose of astronomy: They rarely interact at all with matter. • Thus when a neutrino is formed deep inside a star, it can usually escape to the outside without interacting with any of the matter in the star. • A photon of electromagnetic radiation, on the other hand, can travel only about 0.5 mm (on average) in a stellar interior before it is absorbed, and it takes about a hundred thousand years for a photon to zig and zag its way to the surface.

  33. 12.3c Radioactivity and Neutrinos • The elusiveness of the neutrino not only makes it a valuable messenger—indeed, the only possible direct messenger—carrying news of the conditions inside the Sun at the present time, but also makes it very difficult for us to detect on Earth. • A careful experiment carried out over many years has found only about ⅓ the expected number of neutrinos, as we shall soon see.

  34. 12.4 Stars Shining Brightly • Let us now use our knowledge of atomic nuclei to explain how stars shine. • For a premain-sequence star, the energy from the gravitational contraction goes into giving the individual particles greater speeds; that is, the gas temperature rises. • When atoms collide at high temperature, electrons get knocked away from their nuclei, and the atoms become fully ionized. • The electrons and nuclei can move freely and separately in this “plasma.” • For nuclear fusion to begin, atomic nuclei must get close enough to each other so that the force that holds nuclei together, the “strong nuclear force” (to be discussed in Chapter 19), can play its part. • But all nuclei have positive charges, because they are composed of protons (which bear positive charges) and neutrons (which are neutral). • The positive charges on any two nuclei cause an electrical repulsion between them, which tends to prevent fusion from taking place.

  35. 12.4 Stars Shining Brightly • However, at the high temperatures (millions of kelvins) typical of a stellar interior, some nuclei occasionally have enough energy to overcome this electrical repulsion. • They come sufficiently close to each other that they essentially collide, and the strong nuclear force takes over. • Fusion on the main sequence proceeds in one of two ways, as will be discussed below. • Once nuclear fusion begins, enough energy is generated to maintain the pressure and prevent further contraction. • The pressure provides a force that pushes outward strongly enough to balance gravity’s inward pull.

  36. 12.4 Stars Shining Brightly • In the center of a star, the fusion process is self-regulating. • The star finds a balance between thermal pressure pushing out and gravity pushing in. • It thus achieves stability on the main sequence (at a constant temperature and luminosity). • When we learn how to control fusion in power-generating stations on Earth, which currently seems decades off (and has long seemed so), our energy crisis will be over, since deuterium, the potential “fuel,” is so abundant in Earth’s oceans.

  37. 12.4 Stars Shining Brightly • The greater a star’s mass, the hotter its core becomes before it generates enough pressure to counteract gravity. • The hotter core gives off more energy, so the star becomes brighter (see figure), explaining why main-sequence stars of large mass have high luminosity. • In fact, it turns out that more massive stars use their nuclear fuel at a very much higher rate than less massive stars. • Even though the more massive stars have more fuel to burn, they go through it relatively quickly and live shorter lives than low-mass stars, as we discussed in Chapter 11. • The next two chapters examine the ultimate fates of stars, with the fates differing depending on the masses of the stars.

  38. 12.5 Why Stars Shine • Several chains of nuclear reactions have been proposed to account for the fusion of four hydrogen nuclei into a single helium nucleus. • Hans Bethe of Cornell University suggested some of these procedures during the 1930s. • The different chain reactions prevail at different temperatures, so chains that are dominant in very hot stars may be different from the ones in cooler stars.

  39. 12.5 Why Stars Shine • When the temperature of the center of a main-sequence star is less than about 20 million kelvins, the proton-proton chain (see figure) dominates. • This sequence uses six hydrogen nuclei (protons), and winds up with one helium nucleus plus two protons. • The net transformation is four hydrogen nuclei into one helium nucleus. (Though two of the protons turn into neutrons, here this isn’t the main point.)

  40. 12.5 Why Stars Shine • But the original six protons contained more mass than do the final single helium nucleus plus two protons. • The small fraction of mass that disappears is converted into an amount of energy that we can calculate with the formula E=mc2. • According to Einstein’s special theory of relativity, mass and energy are equivalent and interchangeable, linked by this equation. • For stellar interiors significantly hotter than that of the Sun, the carbon-nitrogen oxygen (CNO) cycle dominates. • This cycle begins with the fusion of a hydrogen nucleus (proton) with a carbon nucleus. • After many steps, and the insertion of four protons, we are left with one helium nucleus plus a carbon nucleus. • Thus, as much carbon remains at the end as there was at the beginning, and the carbon can start the cycle again.

  41. 12.5 Why Stars Shine • As in the proton-proton chain, four hydrogen nuclei have been converted into one helium nucleus during the CNO cycle, 0.7 per cent of the mass has been transformed, and an equivalent amount of energy has been released according to E=mc2. • Main-sequence stars more massive than about 1.1 times the Sun are dominated by the CNO cycle. • Later in their lives, when they are no longer on the main sequence, stars can have even higher interior temperatures, above 108 K. • They then fuse helium nuclei to make carbon nuclei. • The nucleus of a helium atom is called an “alpha particle” for historical reasons. • Since three helium nuclei (2He4) go into making a single carbon nucleus (6C12), the procedure is known as the triple-alpha process.

  42. 12.5 Why Stars Shine • A series of other processes can build still heavier elements inside very massive stars. • These processes, and other element-building methods, are called nucleosynthesis (new´clee-oh-sin´tha-sis). • The theory of nucleosynthesis in stars can account for the abundances (proportions) we observe of the elements heavier than helium. • Currently, we think that the synthesis of isotopes of the lightest elements (hydrogen, helium, and lithium) took place in the first few minutes after the origin of the Universe (Chapter 19), though some of the observed helium was produced later by stars.

  43. 12.6 Brown Dwarfs • When a pre-main-sequence star has at least 7.5 per cent of the Sun’s mass (that is, it has about 75 Jupiter masses), nuclear reactions begin and continue, and it becomes a normal star. • But if the mass is less than 7.5 per cent of the Sun’s mass, the central temperature does not become hot enough for nuclear reactions using ordinary hydrogen (protons) to be sustained. (Masses of this size do, however, fuse deuterium into helium, but this phase of nuclear fusion doesn’t last long because there is so little deuterium in the Universe relative to ordinary hydrogen.)

  44. 12.6 Brown Dwarfs • These objects shine dimly, shrinking and dimming as they age. • They came to be called brown dwarfs, mainly because “brown” is a mixture of many colors and people didn’t agree how such supposedly “failed stars” would look, and also because they emit very little light (see figure). • When old, they have all shrunk to the same radius, about that of the planet Jupiter. • We have met them already in Section 9.2c. • For decades, there was a debate as to whether brown dwarfs exist, but finally some were found in 1995. • We now know of about 1000, because of the advances in astronomical imaging and in spectroscopy, not only in the visible but also in the infrared. • The coolest ones, of spectral type T, show methane and water in their spectra, like giant planets but unlike normal stars.

  45. 12.6 Brown Dwarfs • It is difficult to tell the difference between a brown dwarf and a small, cool, ordinary star, unless the brown dwarf is exceptionally cool. • One way is to see whether an object has lithium in its spectrum. Lithium is a very fragile element, and undergoes fusion in ordinary stars, which converts it to other things. • So if you detect lithium in the spectrum of a dim star, it is probably a brown dwarf (which isn’t sustaining nuclear fusion using protons) rather than a cool, ordinary dwarf star of spectral class M or L, which are the coolest stars on the main sequence (and thus have begun to sustain their nuclear fusion). • A complication is that very young M and L stars might not be old enough to have burned all their lithium, leading to potential confusion with brown dwarfs.

  46. 12.6 Brown Dwarfs • How do we tell the difference between brown dwarfs and giant planets in cases where they are orbiting a more normal star? • Some astronomers would like to distinguish between them by the way that they form: While planets form in disks of dust and gas as the central star is born, brown dwarfs form like the central star, out of the collapse of a cloud of gas and dust. • But we can’t see the history of an object when we look at it, so it is hard to translate the distinction into something observable. • All of the proposed tests are difficult to make. • So, currently, for lack of definitive methods, the distinction is usually made on the basis of mass: Any orbiting object with a mass less than 13 times Jupiter’s is called a planet, while the range 13 to 75 Jupiter masses corresponds to brown dwarfs. (Objects less massive than 13 Jupiter masses not orbiting stars are sometimes called “free-floating planets” since they are not planets in the conventional sense of the word.)

  47. 12.6 Brown Dwarfs • The rationale for using 13 Jupiter masses as the dividing line between planets and brown dwarfs is that above this mass, fusion of deuterium occurs for a short time, whereas below this mass, no fusion ever occurs. • Thus, although brown dwarfs are not normal stars, they do fuse nuclei for a short time, and hence aren’t completely “failed stars” as many people call them. • Brown dwarfs are being increasingly studied, especially in the infrared. • Hubble Space Telescope images show that one of the nearby ones is double, with the components separated by 5 A.U. • By watching it over a few years, we should be able to measure its orbit and derive the masses of the components.

  48. 12.7 The Solar Neutrino Experiment • Astronomers can apply the equations that govern matter and energy in a star, and make a model of the star’s interior in a computer. • Though the resulting model can look quite nice, nonetheless it would be good to confirm it observationally. • However, the interiors of stars lie under opaque layers of gas. • Thus we cannot directly observe electromagnetic radiation from stellar interiors. • Only neutrinos escape directly from stellar cores. • Neutrinos interact so weakly with matter that they are hardly affected by the presence of the rest of the Sun’s mass. • Once formed, they zip right out into space, at (or almost at) the speed of light. • Thus they reach us on Earth about 8 minutes after their birth.

  49. 12.7 The Solar Neutrino Experiment • Neutrinos should be produced in large quantities by the proton-proton chain in the Sun, as a consequence of protons turning into neutrons, positrons, and neutrinos; see the figure. (A positron is an “antielectron,” an example of antimatter. • Whenever a particle and its antiparticle meet, they annihilate each other.)

  50. 12.7a Initial Measurements • For over three decades, astrochemist Raymond Davis has carried out an experiment to search for neutrinos from the solar core, set up in consultation with the theorist John Bahcall, whose calculations long drove the theory. • Davis set up a tank containing 400,000 liters of a chlorine-containing chemical (see figure). • One isotope of chlorine can, on rare occasions, interact with one of the passing neutrinos from the Sun. • It turns into a radioactive form of argon, which Davis and his colleagues at the University of Pennsylvania can detect. • He needs such a large tank because the interactions are so rare for a given chlorine atom. • In fact, he detects fewer than 1 argon atom formed per day, despite the huge size of the tank.

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