Chapter 15--Our Sun - Geological Sciences

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electron neutrinos. However, recent experiments have shown that some of the electron neutrinos might change into muon and tau neutrinos as they fly out through the solar plasma. In that case, our detectors would count fewer than the expected number of electron neutrinos. Early results from the Sudbury Neutrino Observatory in Canada, a new detector designed to search for all types of neutrinos, suggest that neutrinos changing type is indeed the solution to the solar neutrino problem. The observations are ongoing, and it will probably be several more years before this solution can be definitively confirmed. Because of their roles in detecting solar neutrinos and identifying the solar neutrino problem, Raymond Davis, leader of the Homestake experiment, and Masatoshi Koshiba, leader of Super-Kamiokande, shared in the 2002 Nobel Prize for physics. 15.4 From Core to Corona Energy liberated by nuclear fusion in the Sun’s core must eventually reach the solar surface, where it can be radiated into space. The path that the energy takes to the surface is long and complex. In this section, we follow that long path. The Path Through the Solar Interior In Chapter 6, we discussed how atoms can absorb or emit photons. In fact, photons can also interact with any charged particle, and a photon that “collides” with an electron can be deflected into a completely new direction. Deep in the solar interior, the plasma is so dense that the gamma-ray photons resulting from fusion travel only a fraction of a millimeter before colliding with an electron. Because each collision sends the photon in a random new direction, the photon bounces around the core in a haphazard way, sometimes called a random walk.With each random bounce, the photon drifts farther and farther, on average, from its point of origin. As a result, photons from the solar core gradually work their way outward (Figure 15.13). The technical term for this slow, outward migration of photons is radiative diffusion (to diffuse means to “spread out” and radiative refers to the photons of light or radiation). Along the way, the photons exchange energy with their surroundings. Because the surrounding temperature declines as the photons move outward through the Sun, they are gradually transformed from gamma rays to photons of lower energy. (Because energy must be conserved, each gamma-ray photon becomes many lower-energy photons.) By the time the energy of fusion reaches the surface, the photons are primarily visible light. On average, the energy released in a fusion reaction takes about a million years to reach the solar surface. Figure 15.13 A photon in the solar interior bounces randomly among electrons, slowly working its way outward in a process called radiative diffusion. Radiative diffusion is just one type of diffusion. Another is the diffusion of dye through a glass of water. If you place a concentrated spot of dye at one point in the water, each individual dye molecule begins a random walk as it bounces among the water molecules. The result is that the dye gradually spreads through the entire glass. Can you think of any other examples of diffusion in the world around you? Radiative diffusion is the primary way by which energy moves outward through the radiation zone, which stretches from the core to about 70% of the Sun’s radius (see Figure 15.4). Above this point, where the temperature has dropped to about 2 million K, the solar plasma absorbs photons more readily (rather than just bouncing them around). This point is the beginning of the solar convection zone, where the buildup of heat resulting from photon absorption causes bubbles of hot plasma to rise upward in the process known as convection [Section 10.2].Convection occurs because hot gas is less dense than cool gas. Like a hot-air balloon, a hot bubble of solar plasma rises upward through the cooler plasma above it. Meanwhile, cooler plasma from above slides around the rising bubble and sinks to lower layers, where it is heated. The rising of hot plasma and sinking of cool plasma form a cycle that transports energy outward from the top of the radiation zone to the solar surface (Figure 15.14a). The Solar Surface THINK ABOUT IT Earth has a solid crust, so its surface is well defined. In contrast, the Sun is made entirely of gaseous plasma. Defining where the surface of the Sun begins is therefore something like defining the surface of a cloud: From a distance it looks quite distinct, but up close the surface is fuzzy, not sharp. We generally define the solar surface as the layer that appears distinct from a distance. This is the layer we identified as the photosphere when we took our imaginary journey into the Sun. More technically, the photosphere is the layer of the Sun from which photons finally escape into space after the million-year journey of solar energy outward from the core. 508 part V • Stellar Alchemy
a This schematic diagram shows how hot gas (white arrows) rises while cooler gas (orange/black arrows) descends around it. Bright spots appear on the solar surface in places where hot gas is rising from below, creating the granulated appearance of the solar photosphere. Figure 15.14 Convection transports energy outward in the Sun’s convection zone. b Granulation is evident in this VIS photo of the Sun’s surface. Each bright granule is the top of a rising column of gas. At the darker lines between the granules, cooler gas is descending below the photosphere. Each granule is about 1,000 kilometers across. Most of the energy produced by fusion in the solar core ultimately leaves the photosphere as thermal radiation [Section 6.4].The average temperature of the photosphere is about 5,800 K, corresponding to a thermal radiation spectrum that peaks in the green portion of the visible spectrum, with substantial energy coming out in all colors of visible light. The Sun appears whitish when seen from space, but in our sky the Sun appears somewhat more yellow— and even red at sunset—because Earth’s atmosphere scatters blue light. It is this scattered light from the Sun that makes our skies blue [Section 11.3]. Although the average temperature of the photosphere is 5,800 K, actual temperatures vary significantly from place to place. The photosphere is marked throughout by the bubbling pattern of granulation produced by the underlying convection (Figure 15.14b). Each granule appears bright in the center, where hot gas bubbles upward, and dark around the edges, where cool gas descends. If we made a movie of the granulation, we’d see it bubbling rather like a pot of boiling water. Just as bubbles in a pot of boiling water burst on the surface and are replaced by new bubbles, each granule lasts only a few minutes before being replaced by other granules bubbling upward. Sunspots and Magnetic Fields Sunspots are the most striking features on the solar surface (Figure 15.15a). The temperature of the plasma in sunspots is about 4,000 K, significantly cooler than the 5,800 K plasma of the surrounding photosphere. If you think about this for a moment, you may wonder how sunspots can be so much cooler than their surroundings. Why doesn’t the surrounding hot plasma heat the sunspots? Something must be preventing hot plasma from entering the sunspots, and that “something” turns out to be magnetic fields. Detailed observations of the Sun’s spectral lines reveal sunspots to be regions with strong magnetic fields. These magnetic fields can alter the energy levels in atoms and ions and therefore can alter the spectral lines they produce. More specifically, magnetic fields cause some spectral lines to split into two or more closely spaced lines (Figure 15.15b). This effect (called the Zeeman effect) enables scientists to map magnetic fields on the Sun by studying the spectral lines in light from different parts of the solar surface. Magnetic fields are invisible, but in principle we could visualize a magnetic field by laying out many compasses. Each compass needle would point to local magnetic north. We can represent the magnetic field by drawing a series of lines, called magnetic field lines,connecting the needles of these imaginary compasses (Figure 15.16a). The strength of the magnetic field is indicated by the spacing of the lines: Closer lines mean a stronger field (Figure 15.16b). Because these imaginary field lines are so much easier to visualize than the magnetic field itself, we usually discuss magnetic fields by talking about how the field lines would appear. Charged particles, such as the ions and electrons in the solar plasma, follow paths that spiral along the magnetic field lines (Figure 15.16c). Thus, the solar plasma can move freely along magnetic field lines but cannot easily move perpendicular to them. The magnetic field lines act somewhat like elastic bands, being twisted into contortions and knots by turbulent motions in the solar atmosphere. Sunspots occur where the most taut and tightly wound magnetic fields poke nearly straight out from the solar interior. Sunspots tend to occur in pairs connected by a loop of magnetic field lines. These chapter 15 • Our Star 509
Page 1 and 2: PART V STELLAR ALCHEMY
Page 3 and 4: I say Live, Live, because of the Su
Page 5 and 6: Table 15.1 Basic Properties of the
Page 7 and 8: COMMON MISCONCEPTIONS The Sun Is No
Page 9 and 10: in the creation of two gamma-ray ph
Page 11 and 12: Figure 15.9 Vibrations on the surfa
Page 13: Figure 15.11 This tank of dry-clean
Page 17 and 18: weaker magnetic field stronger magn
Page 19 and 20: Figure 15.20 An X-ray image of the
Page 21 and 22: e quator N N N magnetic field line
Page 23 and 24: THE BIG PICTURE Putting Chapter 15
Page 25 and 26: Problems 9. Gravitational Contracti

Chapter 15--Our Sun - Geological Sciences

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