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9-1 The Sun’s energy is generated by thermonuclear reactions inits core

9-2 Energy slowly moves outward from the solar interior throughseveral processes

9-3 The Sun’s outer layers are the photosphere, chromosphere,and corona

9-4 Sunspots are low-temperature regions in thephotosphere

9-5 The Sun’s magnetic field also produces other forms ofsolar activity and causes aurorae on Earth

BY READING THE SECTIONS OF THIS CHAPTER, YOU WILL LEARN

Key Ideas

Our Sun is by far the brightest object in the sky. By earthly standards,

the temperature of its glowing surface is remarkably high, reaching

thousands of degrees. Yet there are regions of the Sun that reach

far higher temperatures of tens of thousands or even millions of degrees.

Gases at such temperatures emit ultraviolet light, which makes them appear

prominent with an ultraviolet telescope in space, as shown in the above image.

Some of the hottest and most energetic regions on the Sun spawn immense

disturbances. These can propel solar material across space far enough to reach

the Earth and other planets.

In recent decades, by looking carefully at the details of how energy is

emitted by the Sun, we have learned that it shines because hundreds of

millions of tons of hydrogen are converted to helium every second

its core. We have also recently come to understand that the Sun has

surprisingly violent atmosphere, with a host of features such as sunspo

whose numbers rise and fall on a predictable 11-year cycle. By studying th

Sun’s vibrations, we have begun to understand new details of the Sun

character far beneath its previously unexplored surface. And, perhaps mo

important, we are in the beginning phases of investigating how changes

the Sun’s activity can affect the Earth’s fragile environment as well as ou

technological society. What we know about the physical processes at wo

inside our closest star helps us understand the stars beyond our sola

system.

P robing the D ynamic Sun

9

R I V U X G A composite view of the Sun showing the upheaval on the surface andthe dynamic outstretched upper atmosphere of the corona. ( SOHO /LASCO/EIT/ESA/NASA)


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210 CHAPTER 9

9-1 The Sun’s energy is generated by thermonuclear reactions in its core

T U

T OR I AL

9 . 1

T

O A9 . 1

If you were to ask the next ve people you meet, “Whatis the most important object in the sky?,” most peoplewould say our Sun. The reasons for the Sun’s importance

are many, including that it provides light to warm Earth’s surface,it provides energy that drives weather, and it underlies the abilityof plants to grow through photosynthesis.

Our Sun also plays an important role in the cosmos. The Sunis the largest member of our solar system. It has almost a thousandtimes more mass than all the solar system’s planets, moons, aster-oids, comets, and meteoroids put together. But the Sun is also a star.In fact, it is a remarkably typical star, with a mass, size, surfacetemperature, and chemical composition that are roughly midwaybetween the extremes exhibited by the myriad of other stars in theheavens.

Solar Energy

For most people, what matters most about the Sun is the energythat it radiates into space. Without the Sun’s warming rays, ouratmosphere and oceans would freeze into an icy layer coating adesperately cold planet, and life on Earth would be impossible. Tounderstand why we are here, we must understand the nature ofthe Sun.

Why is the Sun such an important source of energy? One reasonis that the Sun has a far higher surface temperature than any of theplanets or moons. The Sun’s spectrum is close to that of an idealizedblackbody with a temperature of 5800 K (see Figure 2-12). Thanksto this high temperature, each square meter of the Sun’s surfaceemits a tremendous amount of radiation, principally at visible wave-lengths. Indeed, the Sun is the only object in the solar system thatemits substantial amounts of visible light. The light that we see from

the Moon and planets is actually sunlight that struck those worldsand was reected toward Earth.The Sun’s size also helps us explain its tremendous energy out-

put. Because the Sun is so large, the total number of square metersof radiating surface—that is, its surface area—is immense. Hence,the total amount of energy emitted by the Sun each second, calledits luminosity, is very large indeed: about 3.9 1026 watts, or 3.9 1026 joules of energy emitted every second. Astronomers denotethe Sun’s luminosity by the symbol L

. A circle with a dot in the

center is the astronomical symbol for the Sun and was also used byancient astrologers.

ConceptCheck 9-1 If the Sun emits light at nearly all possiblewavelengths, which range of wavelengths is emitted with the most

intensity?Answer appears at the end of the chapter.

The Source of the Sun’s Energy

What makes the Sun shine so brightly? Albert Einstein discoveredthe underlying key to the energy source within stars in 1905. Ac-cording to his special theory of relativity, a quantity m of mass can

in principle be converted into an amount of energy E according toa now-famous equation:

Einstein’s mass-energy equation

E mc2

E

amount of energy into which the mass can be converted, injoulesm quantity of mass, in kgc speed of light 3 108 m/s

The speed of light c is a large number, so c2 is huge. Therefore,a small amount of matter can release an awesome amount ofenergy.

Einstein didn’t fully appreciate at the time how tremendouslyhis ideas would impact astronomy; it turns out that the tempera-tures and pressures deep within the core of the Sun are so intensethat hydrogen nuclei can combine to produce helium nuclei in anuclear reaction that transforms a tiny amount of mass into a largeamount of energy. This process of converting hydrogen into heliumis called thermonuclear fusion. (It is also sometimes called thermo-

nuclear burning, even though nothing is actually burned in theconventional sense. Ordinary burning involves chemical reactionsthat rearrange the outer electrons of atoms but have no effect onthe atoms’ nuclei.) Thermonuclear fusion can take place only atextremely high temperatures. The reason is that all atomic nucleihave a positive electric charge and so tend to repel one another. Butin the extreme heat and pressure at the Sun’s center, positivelycharged hydrogen nuclei are moving so fast that they can overcometheir electric repulsion and actually touch one another and combine.On Earth, the same thermonuclear fusion provides the devastatingenergy released in a hydrogen bomb.

ANALOGY You can think of protons as tiny electrically chargedspheres that are coated with a very powerful glue. If the spheres are

not touching, the repulsion between their charges pushes themapart. But if the spheres are forced into contact, the strength of theglue “fuses” them together.

CAUTION Be careful not to confuse thermonuclear fusion with thesimilar-sounding process of nuclear ssion. In nuclear fusion, energyis released by joining together nuclei of lightweight atoms such ashydrogen. In nuclear ssion, by contrast, the nuclei of very massiveatoms such as uranium or plutonium release energy by fragmentinginto smaller nuclei. Nuclear power plants produce energy usingssion, not fusion. (Generating power using fusion has been a goalof researchers for decades, but no one has yet devised a commer-cially viable way to do this.)

ConceptCheck

9-2 If hydrogen nuclei are positively charged,under what conditions can two hydrogen nuclei overcome electricalcharge repulsion and combine into helium nuclei, thus releasing energy

according to Einstein’s equation, E mc2?

CalculationCheck 9-1 How much energy is released when just 5kg of mass is converted into energy?

Answers appear at the end of the chapter.


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Probing the Dynamic Sun 21

Converting Hydrogen to Helium

Without its single electron, the nucleus of a hydrogen atom (H) isthe same thing as a single proton. In much the same way, a heliumatom (He) nuclei, in the absence of its two electrons, consists of twoprotons and two neutrons. When they combine, with a concurrentrelease of energy, we can write the nuclear reaction as:

4 H→ He energy

In several separate reactions, two of the four protons arechanged into neutrons, and eventually combine with the remainingprotons to produce a helium nucleus. This sequence of reactionsis called the proton-proton chain (see Cosmic Connections: TheProton-Proton Chain). Each time this process takes place, a smallfraction (0.7%) of the combined mass of the hydrogen nuclei doesnot show up in the mass of the helium nucleus. This “lost” mass isconverted into energy.

CAUTION You may have heard the idea that mass is always con-served (that is, it is neither created nor destroyed), or that energy is

always conserved in a reaction. Einstein’s ideas show that neitherof these statements is quite correct, because mass can be convertedinto energy and vice versa. A more accurate statement is that thetotal amount of mass plus energy is conserved. Hence, the destruc-tion of mass in the Sun does not violate any laws of nature.

For every four hydrogen nuclei converted into a helium nucleus,4.3 1012 joules of energy is released. This may seem like only atiny amount of energy, but it is about 107 times larger than theamount of energy released in a typical chemical reaction, such asoccurs in ordinary burning. To produce the Sun’s luminosity of 3.9 1026 joules per second, 6 1011 kg (600 million metric tons) ofhydrogen must be converted into helium each second. This rate isprodigious, but there is literally an astronomical amount of hydro-gen in the Sun. In particular, the Sun’s core contains enough hydro-gen to have been giving off energy at the present rate for as long asthe solar system has existed, about 4.56 billion years, and to con-tinue doing so for more than 6 billion years into the future.

ConceptCheck 9-3 If 1 kg of hydrogen combines to form heliumin the proton-proton chain, why is only 0.007 kg (0.7%) available to be

converted into energy?

ConceptCheck 9-4 How do astronomers estimate that our Sunhas a lifetime of about 10 billion years?


9-2 Energy slowly moves outward from the solar interior through several processes

T U T O

R IAL 9

. 2

U T O

IA

. 2 While thermonuclear fusion is the source of the Sun’senergy, this process cannot take place everywherewithin the Sun. Extremely high temperatures—in excess

of 107 K—are required for atomic nuclei to fuse together to form

larger nuclei. The temperature of theSun’s visible surface, about 5800 K, is fartoo low for these reactions to occur there.Hence, fusion of atoms releasing energycan be taking place only within the Sun’sinterior. But precisely where does it takeplace? And how does the energy pro-duced by fusion make its way to the sur-face, where it is emitted into space in the form of photons?

To answer these questions, we must understand conditionin the Sun’s interior. Ideally, we would send an exploratory spacecraft to probe deep into the Sun; in practice, the Sun’s intense hewould vaporize even the sturdiest spacecraft. Instead, astronomers use the laws of physics to construct a theoretical model othe Sun.

Hydrostatic and Thermal Equilibrium

Note rst that the Sun is neither growing or shrinking, nor is quickly becoming either hotter or cooler. To understand the naturof this stability, imagine a slab of material in the solar interio

(Figure 9-1a). In equilibrium, the slab on average will move neitheup nor down. (In fact, there are upward and downward motions omaterial inside the Sun, but these motions average out in the lonrun.) Equilibrium is maintained by a balance among three forcethat act on this slab:

1. The downward pressure of the layers of solar material abovthe slab.

2. The upward pressure generated by hot gases beneath the sla

3. The slab’s weight—that is, the downward gravitational pull feels from the rest of the Sun.

The pressure from below must balance both the slab’s weighand the pressure from above. Hence, the pressure below the slamust be greater than that above the slab. In other words, pressur

has to increase with increasing depth. For the same reason, pressurincreases as you dive deeper into the ocean (Figure 9-1b) or as yomove toward lower altitudes in our atmosphere.

In much the same way, we can make inferences about the slabdensity. If it is too dense, its weight will be too great and it will sinkif the density is too low, the slab will rise. To prevent this, the densitof solar material must have a certain value at each depth within thsolar interior. (The same principle applies to objects that oat bneath the surface of the ocean. Scuba divers wear weight belts tincrease their average density so that they will neither rise nor sinbut will stay submerged at the same level.) Astronomers refer tthis equilibrium state of a star, such as the Sun, as being in hydrostatic equilibrium.

Another consideration is that the Sun’s interior is so hot tha

it is completely gaseous. Gases compress and become more denswhen you apply greater pressure to them, so density must increasalong with pressure as you go to greater depths within the SunFurthermore, when you compress a gas, its temperature tends trise, so the temperature must also increase as you move toward thSun’s center.

While the temperature in the solar interior is different at diferent depths, the temperature at each depth remains constant i

Kelvin temperature sca

covered in Box 2-1.

Photons are introduced

Section 2-2.


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212 CHAPTER 9

time. All the energy generated by thermonuclear reactions in theSun’s core must be transported to the Sun’s glowing surface, whereit can be radiated into space. If too much energy owed from thecore to the surface to be radiated away, the Sun’s interior wouldcool down; the Sun’s interior would heat up if too little energyowed to the surface. This principle describing the Sun is calledthermal equilibrium.

ConceptCheck 9-5 If our Sun were much less massive and onlyone-half the diameter, how would the pressure at the Sun’s center be

different from what it actually is?

ConceptCheck 9-6 If our Sun were not in thermalequilibrium and too little energy successfully made it to the surface,

how would the Sun’s core be different?


Transporting Energy Outwardfrom the Sun’s Core

But exactly how is energy transported from the Sun’s center to itssurface? There are three methods of energy transport: conduction,convection, and radiative diffusion. Only the last two are importantinside the Sun.

If you heat one end of a metal bar with a blowtorch, energyows to the other end of the bar so that it too becomes warm. Theef ciency of this method of energy transport, called conduction,varies signicantly from one substance to another. For example,copper is a good conductor of heat, but wood is not (which is whycopper pots often have wooden handles). Conduction is not an ef-cient means of energy transport in substances with low averagedensities, including the gases inside stars like the Sun.

Inside stars like our Sun, energy moves from center to surfaceby two other means: convection and radiative diffusion. Convection is the circulation of uids—gases or liquids—between hot and coolregions. Hot gases rise toward a star’s surface, while cool gases sinkback down toward the star’s center. This physical movement ofgases transports heat energy outward in a star, just as the physicalmovement of water boiling in a pot transports energy from thebottom of the pot (where the heat is applied) to the cooler water atthe surface (see Figure 5-11).

In radiative diffusion, photons emitted from the thermonuclearinferno at a star’s center diffuse outward toward the star’s surface.Individual photons are absorbed and reemitted by atoms and elec-trons inside the star. The overall result is an outward migration fromthe hot core, where photons are constantly created, toward thecooler surface, where they escape into space.

ConceptCheck 9-7 Why is the energy transport process ofconduction relatively unimportant when studying how energy moves

toward the Sun’s surface?

Answer appears at the end of the chapter.

Modeling the Sun

To construct a model of a star like the Sun, astrophysicists express

the ideas of hydrostatic equilibrium, thermal equilibrium, and en-ergy transport as a set of equations. To ensure that the model appliesto the particular star under study, they also make use of astronomi-cal observations of the star’s surface. (For example, to construct amodel of the Sun, they use the data that the Sun’s surface tempera-ture is 5800 K, its luminosity is 3.9 1026 W, and the gas pressureand density at the surface are almost zero.) The astrophysicists thenuse a computer to solve their set of equations and calculate

(a) Material inside the Sun is in hydrostatic equilibrium, so forces balance

Weight of the fish

(b) A fish floating in water is in hydrostatic equilibrium, so forces balance

Pressure from waterbeneath the fish

Pressure from waterabove the fish

Figure 9-1Hydrostatic Equilibrium (a) Material in the Sun’s interior tends to move neither up nor

down. The upward forces on a slab of solar material (due to pressure of gases below theslab) must balance the downward forces (due to the slab’s weight and the pressure of gases

above the slab). Hence, the pressure must increase with increasing depth. (b) The same

principle applies to a fish floating in water. In equilibrium, the forces balance and the fishneither rises nor sinks. (Ken Usami/PhotoDisc)

Pressure from gasesabove the slab

Pressure from gasesbelow the slab

Slab of solarmaterial

Weight of the slab


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The most common form of hydrogen fusion in the Sun involves three

steps, each of which releases energy.

STEP 1

STEP 2

STEP 3

(b) One of the protons changes into a neutron(shown in blue). The proton and neutronform a hydrogen isotope (2H).

(c) One by-product of converting aproton to a neutron is a neutral,nearly massless neutrino ().This escapes from the Sun.

(d) The other by-product of converting a proton to a neutron is apositively charged electron, or positron (e). This encountersan ordinary electron (e), annihilating both particles andconverting them into gamma-ray photons (). The energy ofthese photons goes into sustaining the Sun’s internal heat.

(a) Two protons(hydrogen nuclei,1H) collide.

(a) The 2H nucleusproduced in Step1 collides with athird proton (1H).

(a) The 3He nucleusproduced in Step 2collides withanother 3He nucleusproduced fromthree other protons.

(b) The result of the collision is a helium isotope(3He) with two protons and one neutron.

(b) Two protons and two neutrons from the two3He nuclei rearrange themselves into adifferent helium isotope (4He).

(c) The two remaining protons arereleased. The energy of their motioncontributes to the Sun’s internal heat.

(d) Six 1H nuclei went into producing the two 3He nuclei,which combine to make one 4He nucleus. Since two of theoriginal 1H nuclei are returned to their original state, wecan summarize the three steps as:

4 1H 4He energy

e

e

(c) This nuclear reaction releases another gamma-ray photon(). Its energy also goes into sustaining the internal heat ofthe Sun.

2H

1H

1H

1H

1H

1H

3He

3He

4He

3He

2H

Hydrogen fusion also takes place in all of

the stars visible to the naked eye. (Fusion

follows a different sequence of steps in

the most massive stars, but the net result

is the same.)

Hydrogen fusion in the Sun usually takes

place in a sequence of steps called the

proton-proton chain. Each of these steps

releases energy that heats the Sun and

gives it its luminosity.

The Proton-Proton Chain

(Courtesy of Wally Pacholka)


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214 CHAPTER 9

conditions layer by layer in toward the star’s center. The result is amodel of how temperature, pressure, and density increase with in-creasing depth below the star’s surface.

Table 9-1 and Figure 9-2 show a theoretical model of the Sunthat was calculated in just this way. Different models of the Sun useslightly different assumptions, but all models give essentially thesame results as those shown here. From such computer models wehave learned that at the Sun’s center the density is 160,000 kg/m 3 (14 times the density of lead!), the temperature is 1.55 107 K, and

the pressure is 3.4 1011 atm. (One atmosphere, or 1 atm, is theaverage atmospheric pressure at sea level on Earth.)

Table 9-1 and Figure 9-2 show that the solar luminosity rises to100% at about one-quarter of the way from the Sun’s center to itssurface. In other words, the Sun’s energy production occurs within avolume that extends out only to 0.25 R

. (The symbol R

denotes

the solar radius, or radius of the Sun as a whole, equal to 696,000km.) Outside 0.25 R

, the density and temperature are too low for

thermonuclear reactions to take place. Also note that 94% of the

Distance Pressurefrom the relative toSun’s center Fraction of Fraction Temperature Density pressure(solar radii) luminosity of mass ( 106 K) (kg/m3) at center

0.0 0.00 0.00 15.5 160000 1

0.1 0.42 0.07 13.0 90000 0.46

0.2 0.94 0.35 9.5 40000 0.15

0.3 1.00 0.64 6.7 13000 0.04

0.4 1.00 0.85 4.8 4000 0.007

0.5 1.00 0.94 3.4 1000 0.001

0.6 1.00 0.98 2.2 400 0.0003

0.7 1.00 0.99 1.2 80 4 105

0.8 1.00 1.00 0.7 20 4 106

0.9 1.00 1.00 0.3 2 3 107

1.0 1.00 1.00 0.006 0.00030 4 1013

A Theoretical Model of the SunTABLE 9-1

Distance from Sun’s center (solar radii)


M a s s ( % )

0.2 0.4 0.6 0.8 1.0

L u m i n o s i t y ( % ) 100

75

50

25

100

75

50

25

0.2 0.4 0.6 0.8 1.0

Center of Sun Surface of Sun

D e n s i t y ( k g / m 3 )

160,000

120,000

80,000

40,000

0.2 0.4 0.6



Surface of SunCenter of Sun

0.8 1.0

T e m p e r a t u r e

( 1 0 6 K )

16

12

8

4

0.2 0.4 0.6 0.8 1.0

Surface of SunCenter of Sun

Center of Sun Surface of Sun

I N

T E R A C TIV E

E X .

9 . 1

Figure 9-2A Theoretical Model of the Sun’s Interior These graphs depict what

percentage of the Sun’s total luminosity is produced within each distance from

the center (upper left), what percentage of the total mass lies within each distance

from the center (lower left), the temperature at each distance (upper right), and the

density at each distance (lower right). (See Table 9-1 for a numerical version of this

model.)


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total mass of the Sun is found within the inner 0.5 R. Hence, the

outer 0.5 R contains only a relatively small amount of material.

How energy ows from the Sun’s center toward its surfacedepends on how easily photons move through the gas. If the solargases are comparatively transparent, photons can travel moderatedistances before being scattered or absorbed, and energy is thustransported by radiative diffusion. If the gases are comparativelyopaque, photons cannot get through the gas easily and heat buildsup. Convection then becomes the most ef cient means of energytransport. The gases start to churn, with hot gas moving upwardand cooler gas sinking downward.

From the center of the Sun out to about 0.71 R, energy is

transported by radiative diffusion. Hence, this region is called the ra-diative zone. Beyond about 0.71 R

, the temperature is low enough

(a mere 2 106 K or so) for electrons and hydrogen nuclei to joininto hydrogen atoms. These atoms are very effective at absorbingphotons, much more so than at absorbing free electrons or nuclei,and this absorption chokes off the outward ow of photons. There-fore, beyond about 0.71 R

, radiative diffusion is not an effective

way to transport energy. Instead, convection dominates the energyow in this outer region, which is why it is called the convective

zone. Figure 9-3 shows these aspects of the Sun’s internal structure.Although energy travels through the radiative zone in the form

of photons, the photons have a dif cult time of it. As Table 9-1shows, the material in this zone is extremely dense, so photons fromthe Sun’s core take a long time to diffuse through the radiative zone.As a result, it takes approximately 170,000 years for energy createdat the Sun’s center to travel 696,000 km to the solar surface andnally escape as sunlight. The energy ows outward at an averagerate of 50 centimeters per hour, or about 20 times slower than asnail’s pace.

Once the energy escapes from the Sun, it travels much faster—at the speed of light. Thus, solar energy that reaches you today tooonly 8 minutes to travel the 150 million kilometers from the Sunsurface to the Earth. But this energy was actually produced by themonuclear reactions that took place in the Sun’s core hundreds othousands of years ago.

ConceptCheck 9-8 Which of the following decreases when wemove from the Sun’s central core outward: temperature, mass, or

luminosity?

CalculationCheck 9-2 By what percentage does the Sun’stemperature drop from its central core temperature moving out to a

distance of one-half its radius?


Probing the Sun’s Interior

If the Sun’s interior is not visible from the surface, how might yogo about guring out what is inside? For that matter, how mighyou determine if a melon is ripe at your local grocery store withou

cutting it open? Vibrations are a useful tool for examining the hidden interiors of all kinds of objects. Much like food shoppers whtap melons to listen for particular vibrations and much like geologists who determine the structure of the Earth’s interior by usinseismographs to record vibrations during earthquakes, one poweful technique to infer what is going on beneath the Sun’s surfacinvolves measuring vibrations of the Sun as a whole. This eld osolar research is called helioseismology.

The Sun oscillates in millions of ways as a result of waves resonating in its interior.Figure 9-4 is a computer-generated illustratio

Thermonuclear

energy core

Radiativezone

Convective

zone

0.20.4

0.60.8

Figure 9-3The Sun’s Internal Structure Thermonuclear reactions occur in the Sun’s core, whichextends out to a distance of 0.25 R

from the center. Energy is transported outward, via

radiative diffusion, to a distance of about 0.71 R

. In the outer layers between 0.71 R

and

1.00 R

, energy flows outward by convection.

Convective zoneRadiative zone

Core

Figure 9-4A Sound Wave Resonating in the Sun This computer-generated image shows one of themillions of ways in which the Sun’s interior vibrates. The regions that are moving outward are

colored blue; those moving inward, red. As the cutaway shows, these oscillations are though

to extend into the Sun’s radiative zone (compare Figure 9-3). (National Solar Observatory)


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216 CHAPTER 9

of one such mode of vibration. Helioseismologists can deduce in-formation about the solar interior from measurements of these os-cillations. For example, they have been able to set limits on theamount of helium in the Sun’s core and convective zone and todetermine the thickness of the transition region between the radia-tive zone and convective zone. They have also found that the con-vective zone is thicker than previously thought.

Another approach is to carefully measure everything that comesout of the Sun and then determine how it must have been formed.As part of the process of thermonuclear fusion, protons change intoneutrons and release particles called neutrinos. Like photons, neu-trinos are particles that have no electric charge. Unlike photons,however, neutrinos interact only very weakly with matter. Even thevast bulk of the Sun offers little impediment to their passage, soneutrinos must be streaming out of the core and into space. Indeed,the conversion of hydrogen into helium at the Sun’s center produces1038 neutrinos each second. Every second, about 1014 neutrinoscreated within the Sun must pass through each square meter of theEarth. The challenge is that neutrinos are exceedingly dif cult todetect. Just as neutrinos pass unimpeded through the Sun, they alsopass through the Earth almost as if it were not there. When we are

careful about how we capture them, we are able to conrm that theSun’s energy is indeed caused by thermonuclear reactions just likeour computer models tell us they should.

ConceptCheck 9-9 What can be determined from carefullymonitoring the Sun’s vibrations?


9-3 The Sun’s outer layers are thephotosphere, chromosphere, and corona

Although the Sun’s core is hidden from our direct view, we can easilysee sunlight coming from the high-temperature gases that make upthe Sun’s atmosphere. These outermost layers of the Sun prove tobe the sites of truly dramatic activity, much of which has a directimpact on our planet. By studying these layers, we gain furtherinsight into the character of the Sun as a whole.

Observing the Photosphere

A visible-light photograph like Figure 9-5 makes it appear that theSun has a denite surface. This is actually an illusion; the Sun isgaseous throughout its volume because of its high internal tempera-ture, and the gases simply become less and less dense as you movefarther away from the Sun’s center.

Why, then, does the Sun appear to have a sharp, well-denedsurface? The reason is that essentially all of the Sun’s visible lightemanates from a single, thin layer of gas called the photosphere (“sphere of light”). Just as you can see only a certain distancethrough the Earth’s atmosphere before objects vanish in the haze,we can see only about 400 km into the photosphere. This distanceis so small compared with the Sun’s radius of 696,000 km that thephotosphere appears to be a denite surface. Astronomers usually

dene everything beneath the photosphere as the Sun’s interior andeverything above the photosphere as the Sun’s atmosphere.

Although the photosphere is a very active place, it actually con-tains relatively little material. It has a density of only about 104 kg/ m3, roughly 0.01% the density of the Earth’s atmosphere at sea level.The photosphere is made primarily of hydrogen and helium, themost abundant elements in the solar system. Despite being such athin gas, the photosphere is surprisingly opaque to visible light. If itwere not so opaque, we could see into the Sun’s interior to a depthof hundreds of thousands of kilometers, instead of a mere 400 km.

We can learn still more about the photosphere by examining it

with a telescope—but only when using special dark lters to preventeye damage. Looking directly at the Sun without the correct lter,whether with the naked eye or with a telescope, can cause perma-nent blindness! Under good observing conditions, astronomersusing such lter-equipped telescopes can often see a blotchy patternin the photosphere (Figure 9-6). Each light-colored granule mea-sures about 1000 km (600 mi) across—equal in size to the areas ofTexas and Oklahoma combined—and is surrounded by a darkishboundary. The difference in brightness between the center and theedge of a granule corresponds to a temperature drop of about 300 K.

This granulation appearance is caused by convection of the gasin the photosphere. The inset in Figure 9-6 shows how gas fromlower levels rises upward in granules, cools off, spills over the edgesof the granules, and then plunges back down into the Sun. This can

occur only if the gas is heated from below, like a pot of water beingheated on a stove. Granules form, disappear, and reform in cycleslasting only a few minutes. At any one time, about 4 million gran-ules cover the solar surface.

Superimposed on the pattern of granulation are even largercells, or supergranules, that are about 35,000 km in diameter, largeenough to enclose several hundred granules (Figure 9-7). This large-scale convection moves at only about 0.4 km/s (1400 km/h, or 900

Figure 9-5 R I V U X GThe Photosphere The photosphere is the layer in the solar atmosphere from which theSun’s visible light is emitted. Note that the Sun appears darker around its limb, or edge;

here we are seeing the upper photosphere, which is relatively cool and thus glows less

brightly. (The dark sunspots, which we discuss in Section 9-4, are also relatively cool

regions.) (Celestron International)


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mi/h), about one-tenth the speed of gases churning in a granule, thatcan last about a day.

ANALOGY Similar patterns of large-scale and small-scale convectioncan be found in the Earth’s atmosphere. On the large scale, air risesgradually at a low-pressure area, then sinks gradually at a high-pressure area, which might be hundreds of kilometers away. Thun-derstorms in our atmosphere are small but intense convection cellswithin which air moves rapidly up and down. Like granules, they

last only a relatively short time before they dissipate.

ConceptCheck 9-10 What causes the photosphere to bubblelike water boiling on the stove?


The Sun’s Chromosphere

An ordinary visible-light image such as Figure 9-5 gives the impres-sion that the Sun ends at the top of the photosphere. But during atotal solar eclipse, the Moon blocks the photosphere from our view,revealing a glowing, pinkish layer of gas above the photosphere(Figure 9-8). This is the tenuous chromosphere (“sphere of color”),

the second of the three major levels in the Sun’s atmosphere. Thechromosphere is only about one ten-thousandth (104) as dense asthe photosphere, or about 108 as dense as our own atmosphere.No wonder it is normally invisible!

Unlike the photosphere, which has an absorption line spectrum,the chromosphere has a spectrum dominated by emission lines. Oneof the strongest emission lines in the chromosphere’s spectrum isthe H

a line at 656.3 nm, which is emitted by a hydrogen atom when

Blue: areas of rising gas

Red: areas of sinking gas

Figure 9-7 R I V U X GSupergranules and Large-Scale Convection Supergranules display relatively littlecontrast between their center and edges, so they are hard to observe in ordinary images.

But they can be seen in a false-color Doppler image like this one. Light from gas that is

approaching us (that is, rising) is shifted toward shorter wavelengths, while light from

receding gas (that is, descending) is shifted toward longer wavelengths (see Section 2-5).(David Hathaway, MSFC/NASA)

C o o

l e r

g a s C

o o

l e r

g a s

H o t t e r g a s

H o t t e r g a s

V I D EO 9

. 1 V

E. 1 V I

D EO 9 . 2

O 9

Figure 9-6 R I V U X GSolar Granulation High-resolution photographs of the Sun’s

surface reveal a blotchy pattern called granulation. Granules are convection cells about

1000 km (600 mi) wide in the Sun’s photosphere. The inset shows how rising hot gas

produces bright granules. Cooler gas sinks downward along the boundaries between

granules; this gas glows less brightly, giving the boundaries their dark appearance.This convective motion transports heat from the Sun’s interior outward to the solar

atmosphere. (MSFC/NASA; inset: Goran Scharmer, Lund Observatory)


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218 CHAPTER 9

its single electron falls from the n 3 level to the n 2 level. Thiswavelength is in the red part of the spectrum, which gives the chro-mosphere its characteristic pinkish color. The spectrum also con-tains emission lines of ionized helium. In fact, helium was originallydiscovered in the chromospheric spectrum in 1868, almost 30 yearsbefore helium gas was rst isolated on Earth.

What might be most surprising about the chromosphere is thatthe temperature increases with increasing height in the chromo-sphere. This is just the opposite of the situation in the photosphere,where temperature decreases with increasing height. This is verysurprising, since temperatures should decrease as you move awayfrom the Sun’s interior. In fact, though, the temperature is about4400 K at the top of the photosphere; 2000 km higher, at the topof the chromosphere, the temperature is nearly 25,000 K.

The top photograph in Figure 9-8 is a high-resolution imageof the Sun’s chromosphere taken through an H

a lter. This image

shows numerous vertical spikes, which are actually jets of rising gascalled spicules. A typical spicule lasts just 15 minutes or so: It risesat the rate of about 20 km/s (72,000 km/h, or 45,000 mi/h), canreach a height of several thousand kilometers, and then collapses

and fades away (Figure 9-9). Approximately 300,000 spicules existat any one time, covering about 1% of the Sun’s surface.

Spicules are generally located directly above the edges of gran-ules groups. This is a surprising result, because chromospheric gasesare rising in a spicule while photospheric gases are descending atthe edge of granule groups. What, then, is pulling gases upward toform spicules? The answer proves to be the Sun’s intense magneticeld, discussed in Sections 9-4 and 9-5.

ConceptCheck 9-11 How tall are spicules in the Sun’schromosphere: the height of tall buildings, the distance between large

nearby cities, or the distance across the entire United States?


The Corona

The corona, or outermost region of the Sun’s atmosphere, begins atthe top of the chromosphere. It extends out to a distance of severalmillion kilometers. Despite its tremendous extent, the corona is onlyabout one-millionth (106) as bright as the photosphere—nobrighter than the full moon. Hence, the corona can be viewed onlywhen the light from the photosphere is blocked out, either by useof a specially designed telescope or during a total eclipse.

Figure 9-10 is an exceptionally detailed photograph of theSun’s corona taken during a solar eclipse. It shows that the coronais not merely a spherical shell of gas surrounding the Sun. Rather,

numerous streamers extend in different directions far above thesolar surface. The shapes of these streamers vary on time scalesof days or weeks. The temperatures in the corona are enormousconsidering how far it is from the Sun’s core—temperatures canreach 2 million Kelvins (2 106 K) or even higher—far greaterthan the temperatures in the chromosphere. Figure 9-11 shows howtemperature in both the chromosphere and corona varies withaltitude.

Spicules

Chromosphere

Figure 9-8 R I V U X GThe Chromosphere During a total solar eclipse, the Sun’s glowing chromosphere canbe seen around the edge of the Moon. It appears pinkish because its hot gases emit light

at only certain discrete wavelengths, principally the Ha emission of hydrogen at a red

wavelength of 656.3 nm. The expanded area above shows spicules, jets of chromospheric

gas that surge upward into the Sun’s outer atmosphere. (NOAO)

10,000

8000

6000

4000

2000

D i s t a n c e a b o v e t o p o f p h o t o s p h e r e ( k m )

0

–2000

Photosphere

Chromosphere

Corona

Interior

Spicule Transitionregion

I N T E R

A C TIV E

E X .

9 . 2 Figure 9-9

The Solar Atmosphere This schematic diagram shows the three layers of thesolar atmosphere. The lowest, the photosphere, is about 400 km thick. The chromosphere

extends about 2000 km higher, with spicules jutting up to nearly 10,000 km above the

photosphere. Above a transition region is the Sun’s outermost layer, the corona, which we

discuss in Section 9-3. It extends many millions of kilometers out into space. (Adapted

from J. A. Eddy)


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CAUTION The corona is actually not very “hot”—that is, it containvery little thermal energy. The reason is that the corona is nearly vacuum. In the corona there are only about 1011 atoms per cubmeter, compared with about 1023 atoms per cubic meter in the Sunphotosphere and about 1025 atoms per cubic meter in the air thawe breathe. Because of the corona’s high temperature, the atomthere are moving at very high speeds. But because there are so fewatoms in the corona, the total amount of energy in these movinatoms (a measure of how “hot” the gas is) is rather low. If you ea spaceship into the corona, you would have to worry about becoming overheated by the intense light coming from the photosphere, but you would notice hardly any heating from the coronaultrathin gas.

ANALOGY The situation in the corona is similar to that inside conventional oven that is being used for baking. Both the walls othe oven and the air inside the oven are at the same high temperature, but the air contains very few atoms and thus carries little energy. If you put your hand in the oven momentarily, the lion’s sharof the heat you feel is radiation from the oven walls.

The low density of the corona explains why it is so dim compared with the photosphere. In general, the higher the temperaturof a gas, the brighter it glows. But because there are so few atomin the corona, the net amount of light that it emits is very feeblcompared with the light from the much cooler, but also much densephotosphere.

ConceptCheck 9-12 Why is the corona so difficult to see if it iso much hotter than the photosphere?


The Solar Wind and Coronal Holes

The Earth’s gravity keeps our atmosphere from escaping into spac

In the same way, the Sun’s powerful gravitational attraction keepmost of the gases of the photosphere, chromosphere, and coronfrom escaping. But the corona’s high temperature means that iatoms and ions are moving at very high speeds, around a milliokilometers per hour. As a result, some of the coronal gas can andoes escape. This outow of gas, is called the solar wind.

Each second the Sun ejects about a million tons (10 9 kg) omaterial into the solar wind. But the Sun is so massive that, eveover its entire lifetime, it will eject only a few tenths of a percent oits total mass. The solar wind is composed almost entirely of eletrons and nuclei of hydrogen and helium. About 0.1% of the solawind is made up of ions of more massive atoms, such as siliconsulfur, calcium, chromium, nickel, iron, and argon. The auroraseen at far northern or southern latitudes on Earth are produce

when electrons and ions from the solar wind enter our uppeatmosphere.

Figure 9-12 reveals that the corona is not uniform in temperature or density. The densest, highest-temperature regions appeabright, while the thinner, lower-temperature regions are dark. Notthe large dark area, called a coronal hole because it is almodevoid of luminous gas. Particles streaming away from the Sucan most easily ow outward through these particularly thi

Figure 9-10 R I V U X GThe Solar Corona This striking photograph of the corona was taken during the total solareclipse of July 11, 1991. Numerous streamers extend for millions of kilometers above the

solar surface. The unearthly light of the corona is one of the most extraordinary aspects of

experiencing a solar eclipse. (Courtesy of R. Christen and M. Christen, Astro-Physics, Inc.)

T e m p e r a t u r e ( K )

Height above photosphere (km)

106

105

104

102 103 104 105

Corona

Chromosphere

In this narrow transition region betweenthe chromosphere and corona, the temperaturerises abruptly by about a factor of 100.

Figure 9-11Temperatures in the Sun’s Upper Atmosphere This graph shows how temperature varieswith altitude in the Sun’s chromosphere and corona and in the narrow transition region

between them. In order to show a large range of values, both the vertical and horizontal

scales are nonlinear. (Adapted from A. Gabriel)


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220 CHAPTER 9

Coronal hole

V I D EO

9 . 3

I EO 9 . 3

Figure 9-12 R I V U X GThe Ultraviolet Corona The SOHO spacecraft recorded this false-color

ultraviolet view of the solar corona. The dark feature running across the Sun’s disk

from the bottom is a coronal hole, a region where the coronal gases are thinner than

elsewhere. Such holes are often the source of strong gusts in the solar wind.

( SOHO /EIT/ESA/NASA)

regions. Therefore, it is thought that coronal holes are the maincorridors through which particles of the solar wind escape fromthe Sun.

The temperatures in the corona and the chromosphere are notat all what we would expect. Just as you feel warm if you standclose to a campre but cold if you move away, we would expect

that the temperature in the corona and chromosphere would de-crease with increasing altitude and, hence, increasing distance fromthe warmth of the Sun’s photosphere. Why, then, does the tempera-ture in these regions increase with increasing altitude? This hasbeen one of the major unsolved mysteries in astronomy for the pasthalf-century. As astronomers have tried to resolve this dilemma,they have found important clues in one of the Sun’s most familiarfeatures—sunspots.

ConceptCheck 9-13 From where on the Sun does the solarwind seem to emanate?


9-4 Sunspots are low-temperature regionsin the photosphere

T U

T OR IAL

9 . 3

T U

O IA9 . 3

One might think that the Sun is pretty much the same,day in and day out. Granules, spicules, and the solarwind occur continuously, and these features are said to

be aspects of the quiet Sun. But, as it turns out, other, more dramaticfeatures appear periodically, including massive eruptions and re-gions of concentrated magnetic elds. When these are present, as-tronomers refer to the active Sun. The features of the active Sun thatcan most easily be seen with even a small telescope (although onlywith an appropriate lter attached) are sunspots.

Observing Sunspots

Sunspots are irregularly shaped dark regions in the photosphere.Sometimes sunspots appear in isolation (Figure 9-13a), but fre-quently they are found in sunspot groups (Figure 9-13b; see alsoFigure 9-5). Although sunspots vary greatly in size, typical onesmeasure a few tens of thousands of kilometers across—comparableto the diameter of the Earth. Sunspots are not permanent featuresof the photosphere but last between a few hours and a few months.

Each sunspot has a dark central core, called the umbra, and abrighter border called the penumbra. A sunspot is a region in thephotosphere where the temperature is relatively low, which makesit appear darker than its surroundings. The colors of a sunspotindicate that the temperature of the umbra is typically 4300 K andhat of the penumbra is typically 5000 K. While high by earthlystandards, these temperatures are quite a bit lower than the averagephotospheric temperature of 5800 K. The lower temperature ofsunspots explains why these regions appear so dark.

Occasionally, a sunspot group is large enough to be seen with-out a telescope. Chinese astronomers recorded such sightings 2000years ago, and huge sunspot groups visible to the naked eye (withan appropriate lter) were seen in 1989 and 2003. But it was notuntil Galileo introduced the telescope into astronomy that anyonewas able to examine sunspots in detail. Galileo discovered that hecould determine the Sun’s rotation rate by tracking sunspots asthey moved across the solar disk (Figure 9-14). He found that theSun rotates once in about four weeks. A typical sunspot group lastsabout two months, so a specic one can be followed for two solar

rotations. After more careful study of sunspot movements, it wasdetermined that the equatorial regions rotate more rapidly thanthe polar regions. This phenomenon is known as differential rota-tion. Thus, while a sunspot near the solar equator takes only 25days to go once around the Sun, a sunspot at 30° north or southof the equator takes 27½ days. The rotation period at 75° northor south is about 33 days, while near the poles it may be as longas 35 days.

ConceptCheck 9-14 If the center of a sunspot has atemperature of about 4300 K, why does it appear dark?


The Sunspot CycleThe average number of sunspots on the Sun is not constant, butvaries in a predictable sunspot cycle (Figure 9-15a). This phenom-enon was rst reported by the German astronomer HeinrichSchwabe in 1843 after many years of observing. As Figure 9-15a shows, the average number of sunspots varies with a period of about


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11 years. A period of exceptionally many sunspots is a sunspotmaximum (Figure 9-15b), as occurred in 1979, 1989, and 2000 andprojected to occur in 2013. Conversely, the Sun is almost devoid ofsunspots at a sunspot minimum (Figure 9-15c), as occurred in 1976,1986, 1996, and 2008.

The locations of sunspots also vary with the same 11-year sun-spot cycle. At the beginning of a cycle, just after a sunspot minimum,

sunspots rst appear at latitudes around 30° north and south of thsolar equator (Figure 9-16). Over the succeeding years, the sunspooccur closer and closer to the equator.

Why should the number of sunspots vary with an 11-year cycleWhy should their average latitude vary over the course of a cycleAnd why should sunspots exist at all? The rst step toward answeing these questions came in 1908, when the American astronomeGeorge Ellery Hale discovered that sunspots are associated with intense magnetic elds on the Sun.

When Hale focused a spectroscope on sunlight coming fromsunspot, he found that many spectral lines appear to be split intseveral closely spaced lines (Figure 9-17). This “splitting” of spectrlines is called the Zeeman effect, after the Dutch physicist PieteZeeman, who rst observed it in his laboratory in 1896. Zeemashowed that a spectral line splits when the atoms are subjected tan intense magnetic eld. The more intense the magnetic eld, thwider the separation of the split lines. For more information, se

Looking Deeper 9.1: The Zeeman Effect.Hale’s discovery showed that sunspots are places where the ho

gases of the photosphere are bathed in a concentrated magneteld. Many of the atoms of the Sun’s atmosphere are ionized duto the high temperature. The solar atmosphere is thus a special typof gas called a plasma, in which electrically charged ions and eletrons can move freely. Like any moving, electrically charged objec

November 9

November 12

November 14

November 15

November 17

November 19

V I D EO 9 . 5

. 5

Figure 9-14 R I V U X GTracking the Sun’s Rotation with Sunspots This series of photographs

taken in 1999 shows the rotation of the Sun. By observing the same group of sunspots

from one day to the next, Galileo found that the Sun rotates once in about four weeks.

(The equatorial regions of the Sun actually rotate somewhat faster than the polar

regions.) Notice how the sunspot group shown here changed its shape. (The Carnegie

Observatories)

(a)

Umbra

Penumbra

(b)

V I D EO 9 . 4

O 9 . 4

Figure 9-13 R I V U X GSunspots (a) This high-resolution photograph of the photosphere shows a

mature sunspot. The dark center of the spot is called the umbra. It is bordered

by the penumbra, which is less dark and has a featherlike appearance. (b) In this view

of a typical sunspot group, several sunspots are close enough to overlap. In both

images you can see granulation in the surrounding, undisturbed photosphere. (NOAO)


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222 CHAPTER 9

compass needles all pointing in random directions. However, asHale discovered, however, there is a striking regularity in the mag-netization of sunspot groups. As a given sunspot group moves withthe Sun’s rotation, the sunspots in front are called the “precedingmembers” of the group. The spots that follow behind are referredto as the “following members.” Hale compared the sunspot groupsin the two solar hemispheres, north or south of the Sun’s equator.He found that the preceding members in one solar hemisphere allhave the same magnetic polarity, while the preceding members in

the other hemisphere have the opposite polarity. Furthermore, inthe hemisphere where the Sun has its north magnetic pole, the pre-ceding members of all sunspot groups have north magnetic polarity.In the opposite hemisphere, where the Sun has its south magneticpole, the preceding members all have south magnetic polarity.

they can be deected by magnetic elds. Figure 9-18 shows how amagnetic eld in the laboratory bends a beam of fast-moving elec-trons into a curved trajectory. Similarly, the paths of moving ionsand electrons in the photosphere are deected by the Sun’s magneticeld. In particular, magnetic forces act on the hot plasma that risesfrom the Sun’s interior due to convection. Where the magnetic eldis particularly strong, these forces push the hot plasma away. Theresult is a localized region where the gas is relatively cool and thusglows less brightly—in other words, a sunspot. By carefully measur-

ing the magnetic elds around a sunspot group, we discover that agroup resembles a giant bar magnet, with a north magnetic pole atone end and a south magnetic pole at the other.

If different sunspot groups were unrelated to one another,their magnetic poles would be randomly oriented, like a bunch of

S o l a r l a t i t u d e

1880 1890 1900 1910 1920 1930 1940 1950 1960 1970 1980 1990 2000 2010

Date

90 N

30 N

0

30 S

90 S

Figure 9-16Variations in the Average Latitude of Sunspots The dots in this graph (sometimes calleda “butterfly diagram”) record how far north or south of the Sun’s equator sunspots were

observed. At the beginning of each sunspot cycle, most sunspots are found near latitudes

30° north or south. As the cycle goes on, sunspots typically form closer to the equator.

(NASA Marshall Space Flight Center)

300

200

100

0 A v e r a g e n u m b e r

o f s u n s p o t s

1750 1770 1790 1810 1830 1850 1870 1890 1910 1930 1950

Date

1970 1990 2010

(b) Near sunspot maximum (c) Near sunspot minimum

Figure 9-15 R I V U X GThe Sunspot Cycle (a) The number of sunspots on the Sunvaries with a period of about 11 years. The most recent sunspot

maximum occurred in 2000. (b) This photograph, taken near

sunspot maximum in 1989, shows a number of sunspots and

large sunspot groups. The sunspot group visible near the bottomof the Sun’s disk has about the same diameter as the planet

Jupiter. (c) Near sunspot minimum, as in this 1986 photograph,

essentially no sunspots are visible. (NOAO)


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Along with his colleague Seth B. Nicholson, Hale also discovered that the Sun’s polarity pattern completely reverses itself ever11 years—the same interval as the time from one solar maximumto the next. The hemisphere that has preceding north magnetipoles during one 11-year sunspot cycle will have preceding soutmagnetic poles during the next 11-year cycle, and vice versa. Thnorth and south magnetic poles of the Sun itself also reverse ever11 years. Thus, the Sun’s magnetic pattern repeats itself only aftetwo sunspot cycles, which is why astronomers speak of a 22-yeasolar cycle.

ConceptCheck 9-15 Is the sunspot cycle an 11-year cycle or a22-year cycle?


The Magnetic-Dynamo Model

In 1960, the American astronomer Horace Babcock proposed

description that seems to account for many features of this 22-yeasolar cycle. Babcock’s scenario, called a magnetic-dynamo modemakes use of two basic properties of the Sun’s photosphere—differential rotation and convection. Differential rotation causes thmagnetic eld in the photosphere to become wrapped around thSun (Figure 9-19). As a result, the magnetic eld becomes concentrated at certain latitudes on either side of the solar equator. Convection in the photosphere creates tangles in the concentrate

(b) The spectrum in and around the sunspot

(a) A sunspot

Outside the sunspot, the

magnetic field is lowand this iron absorptionline is single.

Within the sunspot, themagnetic field is strongand this iron absorptionline splits into three.

Figure 9-17 R I V U X GSunspots Have Strong Magnetic Fields (a) A black line in this image of a sunspot showswhere the slit of a spectrograph was aimed. (b) This is a portion of the resulting spectrum,

including a dark absorption line caused by iron atoms in the photosphere. The splitting of

this line by the sunspot’s magnetic field can be used to calculate the field strength. Typica

sunspot magnetic fields are over 5000 times stronger than the Earth’s field at its north and

south poles. (NOAO)

Figure 9-18 R I V U X GMagnetic Fields Deflect Moving, Electrically Charged Objects In this laboratoryexperiment, a beam of negatively charged electrons (shown by a blue arc) is aimed

straight upward from the center of the apparatus. The entire apparatus is inside a large

magnet, and the magnetic field deflects the beam into a curved path. (Courtesy of

Central Scientific Company)


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224 CHAPTER 9

magnetic eld, and “kinks” erupt through the solar surface. Sun-spots appear where the magnetic eld protrudes through the pho-tosphere. The theory suggests that sunspots should appear rst atnorthern and southern latitudes and later form nearer to the equa-tor. This is just what is observed (see Figure 9-16). Note also that,as shown on the far right in Figure 9-19, the preceding member ofa sunspot group has the same polarity (N or S) as the Sun’s magneticpole in that hemisphere.

Differential rotation eventually undoes the twisted magneticeld. The preceding members of sunspot groups move toward theSun’s equator, while the following members migrate toward thepoles. Because the preceding members from the two hemisphereshave opposite magnetic polarities, their magnetic elds cancel eachother out when they meet at the equator. The following membersin each hemisphere have the opposite polarity to the Sun’s polein that hemisphere; hence, when they converge on the pole, the fol-lowing members rst cancel out and then reverse the Sun’s overallmagnetic eld. The elds are now completely relaxed. Once again,differential rotation begins to twist the Sun’s magnetic eld, butnow with all magnetic polarities reversed. In this way, Babcock’smodel helps to explain the change in eld direction every 11 years.

By comparing the speeds of sound waves that travel with and

against the Sun’s rotation, astronomers now understand that theSun’s rotation rate is different at different depths and latitudes. Asshown in Figure 9-20, the Sun’s surface pattern of differential rota-tion persists through the convective zone. Farther in, within theradiative zone, the Sun seems to rotate like a rigid object with aperiod of 27 days at all latitudes. Astronomers suspect that the Sun’smagneticeld originates in a relatively thin layer where the radiative

Start

N

After 1 rotation After 2 rotations After 3 rotationsAftermany

rotations

S N

N S

S

Start

N

After 1 rotation After 2 rotations After 3 rotationsAftermany

rotations

S N

N S

S

Figure 9-19Babcock’s Magnetic-Dynamo Model Magnetic field lines tend to move along with theplasma in the Sun’s outer layers. Because the Sun rotates faster at the equator than near the

poles, a field line that starts off running from the Sun’s north magnetic pole (N) to its south

magnetic pole (S) ends up wrapped around the Sun like twine wrapped around a ball. The

insets on the far right show how sunspot groups appear where the concentrated magnetic

field rises through the photosphere.

25 Days 35 Days

Figure 9-20Rotation of the Solar Interior This cutaway picture of the Sun shows how thesolar rotation period (shown by different colors) varies with depth and latitude. The

surface and the convective zone have differential rotation (a short period at the

equator and longer periods near the poles). Deeper within the Sun, the radiative

zone seems to rotate like a rigid sphere. (Courtesy of K. Libbrecht, Big Bear Solar

Observatory)


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Magnetic Arches

In a plasma, magnetic eld lines and the material of the plasma tento move together. The tendency of plasma to follow the Sun’s magnetic eld helps to explain why the temperature of the chromosphere and corona is so high. Spacecraft observations show magneteld arches extending tens of thousands of kilometers into the corona, with streamers of electrically charged particles moving aloneach arch (Figure 9-21a). If the magnetic elds of two arches cominto proximity, their magnetic elds can rearrange and combinThe tremendous amount of energy stored in the magnetic eld then released into the solar atmosphere. (A single arch contains amuch energy as a hydroelectric power plant would generate in million years.) The amount of energy released in this way appearto be more than enough to maintain the temperatures of the chromosphere and corona.

ANALOGYThe idea that a magnetic eld can heat gases has applications on Earth as well as on the Sun. In an automobile engine’s igntion system an electric current is set up in a coil of wire, whicproduces a magnetic eld. When the current is shut off, the magnet

eld collapses and its energy is directed to a spark plug in one othe engine’s cylinders. The released energy heats the mixture of aand gasoline around the plug, causing the mixture to ignite. Thdrives the piston in that cylinder and makes the automobile go.

V I D EO 9 . 6

D 9 . 6 Magnetic heating can also explain why the parts of thcorona that lie on top of sunspots are often the mosprominent in ultraviolet images. (Some examples are th

bright regions in Figure 9-12.) The intense magnetic eld of th

and convective zones meet and slide past each other due to theirdifferent rotation rates.

Adding to the yet-to-be-fully-understood nature of the Sun,there seem to be times when all traces of sunspots and the sunspotcycle vanish for many years. For example, virtually no sunspots wereseen from 1645 through 1715. Curiously, during these same yearsEurope experienced record low temperatures, often referred to asthe Little Ice Age, whereas the western United States was subjectedto severe drought. By contrast, there was apparently a period of in-creased sunspot activity during the eleventh and twelfth centuries,during which the Earth was warmer than it is today. Thus, variationsin solar activity appear to affect climates on the Earth. The origin ofthis Sun-Earth connection is a topic of ongoing research.

ConceptCheck 9-16 How might the Sun’s sunspot cycle changeif the Sun were rotating much faster than it is now?


9-5 The Sun’s magnetic field also producesother forms of solar activity and causesaurorae on Earth

If magnetic elds are so powerful on the Sun, what other effectsmight the Sun’s intense magnetic eld be able to cause? In fact, theSun’s magnetic eld does more than just explain the presence ofsunspots.

Coronal loops

Image of Earth

(superimposed for scale)

(a)

Figure 9-21 R I V U X GMagnetic Arches and Magnetic Reconnection (a) This false-color ultraviolet imagefrom the TRACE spacecraft (Transition Region and Coronal Explorer ) shows magnetic fieldloops suspended high above the solar surface. The loops are made visible by the glowing

gases trapped within them. (b) When the magnetic fields in these loops change their

arrangement, a tremendous amount of energy is released and solar material can be ejecte

upward. (a: Stanford-Lockheed Institute for Space Research; TRACE; and NASA)

1. If magnetic

field loopsbegin to pinchtogether . . .

2. . . . the field lines

of adjacent loops canreconnect, causinga release of energy.

3. The upper helix or

“coil” of magnetic fieldcan break loose, carryingmaterial with it into space

(b)


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226 CHAPTER 9

sunspots helps trap and compress hot coronal gas, giving it such ahigh temperature that it emits copious amounts of high-energy ul-traviolet photons and even more energetic X-ray photons.

ConceptCheck 9-17 Why does glowing plasma on the Sunappear to arch up above the Sun’s photosphere?


Prominences, Solar Flares, andCoronal Mass Ejections

Coronal heating occurs even when the Sun is quiet. But magneticelds can also push upward from the Sun’s interior, compress-ing and heating a portion of the chromosphere that appears asbright, arching columns of gas called prominences (Figure 9-22).These can extend for tens of thousands of kilometers above thephotosphere. Some prominences last for only a few hours, whileothers persist for many months. The most energetic prominencesbreak free of the magnetic elds that conned them and burst intospace.

V I D EO

9 .

7

E 9 7 Violent, eruptive events on the Sun, called solar ares,occur in complex sunspot groups. Within only a few min-utes, temperatures in a compact region may soar to 5

106 K, and vast quantities of particles and radiation—including asmuch material as is in the prominence shown in Figure 9-22—areblasted out into space. These eruptions can also cause disturbancesthat spread outward in the solar atmosphere, like the ripples thatappear when you drop a rock into a pond.

Prominence

Bright areaslie on top of

sunspot groups

Figure 9-22 R I V U X GA Solar Prominence A huge prominence arches above the solar surface in this ultravioletimage from the SOHO spacecraft. The image was recorded using light at a wavelength of

30.4 nm, emitted by singly ionized helium atoms at a temperature of about 60,000 K. By

comparison, the material within the arches in Figure 9-21 reaches temperatures in excess of

2 106K. ( SOHO /EIT/ESA/NASA)

Material ejectedfrom the corona

Ejected material encountersEarth’s magnetosphere

Earth

(a) A coronal mass ejection (b) Two to four days later

Figure 9-23 R I V U X GA Coronal Mass Ejection (a) SOHO recorded this coronalmass ejection in an X-ray image. (The image of the Sun itself

was made at ultraviolet wavelengths.) (b) Within two to four

days the fastest-moving ejected material reaches a distance of

1 AU from the Sun. Most particles are deflected by the Earth’s

magnetosphere, but some are able to reach the Earth. (The

ejection shown in (a) was not aimed toward the Earth and did

not affect us.) ( SOHO /EIT/LASCO/ESA/NASA)


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even when the Sun is at its quietest. Astronomers are devoting substantial effort to understanding these and other aspects of our dynamic Sun. Table 9-2 lists essential data about the Sun.

ConceptCheck 9-18 Which of the following are the mostenergetic: prominences, solar flares, or coronal mass ejections?


Key Ideas and Terms

9-1 The Sun’s energy is generated by thermonuclear reactions inits core

• The Sun’s luminosity is the amount of energy emitted each second anis produced by the proton-proton chain in which four hydrogen nuclecombine to produce a single helium nucleus.

• The energy released in a nuclear reaction corresponds to a slightreduction of mass, as predicted by Einstein’s equation E mc2.

• Thermonuclear fusion occurs only at very high temperatures;for example, hydrogen fusion occurs only at temperatures inexcess of about 107 K. In the Sun, fusion occurs only in the dense,hot core.

9-2 Energy slowly moves outward from the solar interior through severaprocesses

The most energetic ares carry as much as 1030 joules of energy,equivalent to 1014 one-megaton nuclear weapons being exploded atonce! However, the energy of a solar are does not come from ther-monuclear fusion in the solar atmosphere; instead, it appears to bereleased from the intense magnetic eld around a sunspot group.

As energetic as solar ares are, they are dwarfed by coronalmass ejections. One such event is shown in the image that opensthis chapter; Figure 9-23a shows another. In a coronal mass ejection,more than 1012 kilograms (a billion tons) of high-temperature coro-nal gas is blasted into space at speeds of hundreds of kilometers persecond. A typical coronal mass ejection lasts a few hours. Theseexplosive events seem to be related to large-scale alterations in theSun’s magnetic eld, like the magnetic reconnection shown in Figure9-21b. Coronal mass ejections occur every few months; smallereruptions may occur almost daily.

If a solar are or coronal mass ejection happens to be aimedtoward Earth, a stream of high-energy electrons and nuclei reachesus a few days later (Figure 9-23b). When this plasma arrives, it caninterfere with satellites, pose a health hazard to astronauts in orbit,

and disrupt electrical and communications equipment on the Earth’ssurface. Telescopes on Earth and on board spacecraft now monitorthe Sun continuously to provide warnings of dangerous levels ofsolar particles.

The numbers of sunspots, prominences, solar ares, and coro-nal mass ejections all vary with the same 11-year cycle as sunspots.But unlike sunspots, coronal mass ejections never completely cease,

Distance from Earth: Mean: 1 AU 149,598,000 km Maximum: 152,000,000 km

Minimum: 147,000,000 km

Light travel time to Earth: 8.32 min Mean angular diameter: 32 arcmin

Radius: 696,000 km 109 Earth radii

Mass: 1.9891 1030 kg 3.33 105 Earth masses

Composition (by mass): 74% hydrogen, 25% helium,

1% other elements

Composition (by number of atoms): 92.1% hydrogen, 7.8% helium,

0.1% other elements

Mean density: 1410 kg/m3

Mean temperatures: Surface: 5800 K; Center: 1.55 107 K

Luminosity: 3.90 1026 W

Distance from center of Galaxy: 8000 pc

26,000 ly

Orbital period around center 220 million years

Orbital speed around center 220 km/sof Galaxy:

Sun DataTABLE 9-2


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228 CHAPTER 9

• Neutrinos emitted in thermonuclear reactions in the Sun’s core havebeen detected, but in smaller numbers than expected. Recent neutrinoexperiments explain why this is so.

• Helioseismology is the study of how the Sun vibrates, which has beenused to infer pressures, densities, chemical compositions, and rotationrates within the Sun.

9-3 The Sun’s outer layers are the photosphere, chromosphere, andcorona

• The visible surface of the Sun, the photosphere, is the lowest layer inthe solar atmosphere. Its spectrum is similar to that of a blackbody ata temperature of 5800 K. Convection in the photosphere produces

granules.

• A theoretical description of a star’s interior can be modeled using thelaws of physics showing that it is in hydrostatic equilibrium whereenergy moving outward precisely balances its gravitational pullinward.

• The standard model of the Sun suggests that hydrogen fusion takesplace in a core extending from the Sun’s center to about 0.25 solar

radius and that our Sun is in thermal equilibrium.

• The core is surrounded by a radiative zone extending to about 0.71 solarradius. In this zone, energy travels outward through radiative diffusion.

• The radiative zone is surrounded by a rather opaque convective zone of gas at relatively low temperature and pressure. In this zone, energytravels outward primarily through convection.

The Sun

PROMPT: What would you tell a fellow student who said, “At the

halfway point between the Sun’s center and its photosphere,it has half the temperature and density of the core, contains half

the Sun’s total mass, and produces half of the Sun’s luminosity.”

ENTER RESPONSE:

Guiding Questions

1. At 0.5 of the Sun’s radius, the temperature is about

a. one-fourth of the core temperature.

b. one-half of the core temperature.

c. the same as the temperature throughout.

d. the same temperature as the photosphere.

2. At 0.5 of the Sun’s radius, the density is about

a. one-third of the core density.

b. one-half of the core density.c. the same density as the photosphere.

d. the same as water.

3. The percentage of mass contained within 0.5 of the Sun’s radiusis about

a. 90%.

b. 50%.

c. 33%.

d. 10%.

4. Nearly all of the Sun’s luminosity is generated within the inner

a. one-third of the radius.

b. one-half of the radius.

c. 0.8 of the radius.

d. 0.2 of the radius.

M a s s ( % )

0.2 0.4 0.6 0.8 1.0

L u m i n o s i t y ( % ) 100

75

50

25

100

75

50

25

0.2 0.4 0.6 0.8 1.0

Center Surface

D e n s i t y ( k g / m 3 )

160,000

120,000

80,000

40,000

0.2 0.4 0.6

Distance from Sun’s center, R.

0.8 1.0

T e m p e r a t u r e

( 1

0 6 K )

16

12

8

4

0.2 0.4 0.6 0.8 1.0


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7. Briey describe the three layers that make up the Sun’s atmosphere.In what ways do they differ from each other?

T U T O

R IAL 9

. 3 8. How do astronomers know when the next sunspot

maximum and sunspot minimum will occur?9. Why do astronomers say that the solar cycle is

really 22 years long, even though the number of sunspots varies

over an 11-year period?

10. Explain how the magnetic-dynamo model accounts for the solar

cycle.

11. Why should solar ares and coronal mass ejections be a concern forbusinesses that use telecommunication satellites?

Web Chat Questions

1. Discuss the extent to which cultures around the world haveworshiped the Sun as a deity throughout history. Why do you

suppose there has been such widespread veneration?

2. In the movie Star Trek IV: The Voyage Home, the starship Enterprisies on a trajectory that passes close to the Sun’s surface. Whatfeatures should a real spaceship have to survive such a ight? Why?

3. Discuss some of the dif culties in correlating solar activity with

changes in Earth’s climate.

4. Describe some of the advantages and disadvantages of observing theSun (a) from space and (b) from Earth’s south pole. What kinds ofphenomena and issues might solar astronomers want to explore fromthese locations?

Collaborative Exercises 1. Figure 9-16 shows variations in the average latitude of sunspots.

Estimate the average latitude of sunspots in the year you were bornand estimate the average latitude on your twenty-rst birthday. Makerough sketches of the Sun during those years to illustrate your answer

2. Create a diagram showing a sketch of how limb darkening on theSun would look different if the Sun had either a thicker or thinnerphotosphere. Be sure to include a caption explaining your diagram.

3. Solar granules, shown in Figure 9-6, are about 1000 km across. Whacity is about that distance away from where you are right now? Whcity is that distance from the birthplace of each group member?

4. Magnetic arches in the corona are shown in Figure 9-21a. Howmany Earths high are these arches, and how many Earths could tinside one arch?

Observing Projects 1. Use the Starry Night College™ program to measure the Sun’s

rotation. Select Favourites > Investigating Astronomy > SolarRotation to display the Sun as seen from about 0.008 AU above itssurface, well inside the orbit of Mercury. Use the Time Flow controlto stop the Sun’s rotation at a time when a line of longitude on theSun makes a straight line between the solar poles, preferably a linecrossing a recognizable solar feature. Note the date and time. Run

time forward and adjust the date and time to place the selectedmeridian in this position again.

a) What is the rotation rate of the Sun as shown in Starry NightCollege™?

b) The demonstration in part (a) does not show one importantfeature of the Sun, namely its differential rotation, in which theequator of this uid body rotates faster than the polar regions.

• Above the photosphere is a layer of less dense but higher temperaturegases called the chromosphere. Spicules extend upward from thephotosphere into the chromosphere.

• The outermost layer of the solar atmosphere, the corona, is made ofvery high-temperature gases at extremely low density. A stream ofparticles making a solar wind emanates from thin regions calledcoronal holes.

9-4 Sunspots are low-temperature regions in the photosphere

• Sunspots are relatively cool regions produced by local concentrationsof the Sun’s magnetic eld.

• The average number of sunspots increases to a sunspot maximum anddecreases to a sunspot minimum in a regular sunspot cycle ofapproximately 11 years, with reversed magnetic polarities from one11-year cycle to the next. Two such cycles make up a 22-year solarcycle in which the surface magnetic eld increases, decreases, and thenincreases again with the opposite polarity.

• The magnetic polarity is measured by observing the Zeemaneffect.

• The magnetic-dynamo model suggests that many features of the solarcycle are due to changes in the Sun’s magnetic eld. These changes are

caused by convection and the Sun’s differential rotation.9-5 The Sun’s magnetic field also produces other forms of solar activityand causes aurorae on Earth

• Plasma on the Sun arranges itself into various observable features,called prominences.

• A solar are is a brief eruption of hot, ionized gases from a sunspotgroup. A coronal mass ejection is a much larger eruption that involvesimmense amounts of gas from the corona.

• When charged particles emitted by the Sun interact with Earth’satmosphere, it causes an aurora where the upper atmosphere glows.When observed in the northern hemisphere it is called the northernlights or aurora borealis.

QuestionsReview Questions 1. What is meant by the luminosity of the Sun?

T U T O

R IAL 9

. 1

2. What is thermonuclear fusion? Why is this fusionfundamentally unlike the burning of a log in areplace?

3. Why do thermonuclear reactions occur only in the Sun’s core, notin its outer regions?

T U T O R

IAL 9

. 2 4. If thermonuclear fusion in the Sun were suddenly to

stop, what would eventually happen to the overallradius of the Sun? Justify your answer using the ideas

of hydrostatic equilibrium and thermal equilibrium. 5. Give some everyday examples of conduction, convection, and

radiative diffusion.

6. What is a neutrino? Why is it useful to study neutrinos coming fromthe Sun? What do they tell us that cannot be learned from otheravenues of research?


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230 CHAPTER 9

a) What is the distance from the Sun of the following stars: RigelKentaurus; Sirius; Fomalhaut; Vega; Arcturus?

b) Which star within the above group has the highesttemperature?

c) Which is intrinsically the most luminous of these stars?

Answers

ConceptChecksConceptCheck 9-1: The Sun emits most of its energy in the form of visiblelight.

ConceptCheck 9-2: At the extremely high temperatures and pressures exist-ing in the Sun’s core, hydrogen nuclei can move fast enough to overcomethe electrical charge repulsion and fuse together into helium nuclei.

ConceptCheck 9-3: When 1 kg of hydrogen combines to form helium, thevast majority of the mass is used as the substance of helium atoms, withonly 0.7% of the original mass left over to be converted into energy.

ConceptCheck 9-4: Astronomers use the current energy output of the Sun

to estimate how fast the Sun is consuming its usable fuel and estimate howmuch fuel it has available to continue at its present consumption rate.

ConceptCheck 9-5: Because pressure in the Sun’s core is due to the down-ward pushing weight of the overlying mass of material, having less mass

pressing down would result in a lower pressure at the core.

ConceptCheck 9-6: When too little energy ows to the surface, the Sun’score temperature would increase dramatically.

ConceptCheck 9-7: The energy transport process of conduction occurswhen energy moves through a rel

the day the earth was as still as darwin

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