Chapter 15--Our Sun - Geological Sciences

PART 

V 

STELLAR 

ALCHEMY

15 Our Star 

LEARNING GOALS 

15.1 Why Does the Sun Shine? 

• What process creates energy in the Sun? 

• Why does the Sun’s size remain stable? 

• How did the Sun become hot enough for fusion 

in the first place? 

15.2 Plunging to the Center of the Sun: 

An Imaginary Journey 

• What are the major layers of the Sun, from the 

center out? 

• What do we mean by the “surface” of the Sun? 

• What is the Sun made of? 

15.3 The Cosmic Crucible 

• Why does fusion occur in the Sun’s core? 

• Why is energy produced in the Sun at such 

a steady rate? 

• Why was the Sun dimmer in the distant past? 

• How do we know what is happening inside the Sun? 

• What is the solar neutrino problem? Is it solved ? 

15.4 From Core to Corona 

• How long ago did fusion generate the energy we 

now receive as sunlight? 

• How are sunspots, prominences, and flares related 

to magnetic fields? 

• What is surprising about the temperature of the 

chromosphere and corona, and how do we explain it? 

15.5 Solar Weather and Climate 

• What is the sunspot cycle? 

• What effect does solar activity have on Earth and 

its inhabitants? 

496

I say Live, Live, because of the Sun, 

The dream, the excitable gift. 

Anne Sexton (1928–1974) 

Astronomy today involves the study of the 

entire universe, but the root of the word 

astronomy comes from the Greek word 

for “star.” Although we have learned a lot about the 

universe up to this point in the book, only now do 

we turn our attention to the study of the stars, the 

namesakes of astronomy. 

When we think of stars, we usually think of the 

beautiful points of light visible on a clear night. The 

nearest and most easily studied star is visible only 

in the daytime—our Sun. Of course, the Sun is important 

to us in many more ways than as an object 

for astronomical study. The Sun is the source of virtually 

all light, heat, and energy reaching Earth, and life 

on Earth’s surface could not survive without it. 

In this chapter, we will study the Sun in some 

depth. We will learn how the Sun makes life possible 

on Earth. Equally important, we will study our Sun as 

a star so that in subsequent chapters we can more 

easily understand stars throughout the universe. 


Ancient peoples recognized the vital role of the Sun in their 

lives. Some worshiped the Sun as a god, and others created 

elaborate mythologies to explain its daily rise and set. Only 

recently, however, have we learned how the Sun provides us 

with light and heat. 

Most ancient thinkers viewed the Sun as some type of 

fire, perhaps a lump of burning coal or wood. The Greek 

philosopher Anaxagoras (c. 500–428 B.C.) imagined the Sun 

to be a very hot, glowing rock about the size of the Greek 

peninsula of Peloponnesus (comparable in size to Massachusetts). 

Thus, he was one of the first people in history to 

believe that the heavens and Earth are made from the same 

types of materials. 

By the mid-1800s, the size and distance of the Sun 

were reasonably well known, and scientists seriously began 

to address the question of how the Sun shines. Two early 

ideas held either that the Sun was a cooling ember that had 

once been much hotter or that the Sun generated energy 

from some type of chemical burning similar to the burning 

of coal or wood. Simple calculations showed that a cooling 

or chemically burning Sun could shine for a few thousand 

years—an age that squared well with biblically based estimates 

of Earth’s age that were popular at the time. However, 

these ideas suffered from fatal flaws. If the Sun were a cooling 

ember, it would have been much hotter just a few hundred 

years earlier, making it too hot for civilization to have 

existed. Chemical burning was ruled out because it cannot 

generate enough energy to account for the rate of radiation 

observed from the Sun’s surface. 

A more plausible hypothesis of the late 1800s suggested 

that the Sun generates energy by contracting in size, 

a process called gravitational contraction.Ifthe Sun were 

shrinking, it would constantly be converting gravitational 

potential energy into thermal energy, thereby keeping the 

Sun hot. Because of its large mass, the Sun would need to 

contract only very slightly each year to maintain its temperature—so 

slightly that the contraction would be unnoticeable. 

Calculations showed that the Sun could shine 

for up to about 25 million years generating energy by gravitational 

contraction. However, geologists of the late 1800s 

had already established the age of Earth to be far older than 

25 million years, leaving astronomers in an embarrassing 

position. 

Only after Einstein published his special theory of 

relativity, which included his discovery of E mc 2 ,did 

the true energy-generation mechanism of the Sun become 

clear. We now know that the Sun generates energy by nuclear 

fusion, a source so efficient that the Sun can shine for 

about 10 billion years. Because the Sun is only 4.6 billion 

years old today [Section 9.5],we expect it to keep shining 

for some 5 billion more years. 

According to our current model of solar-energy generation 

by nuclear fusion, the Sun maintains its size through 

a balance between two competing forces: gravity pulling 

inward and pressure pushing outward. This balance is called 

gravitational equilibrium (or hydrostatic equilibrium). 

It means that, at any point within the Sun, the weight of 

overlying material is supported by the underlying pressure. 

A stack of acrobats provides a simple example of this balance 

(Figure 15.1). The bottom person supports the weight 

of everybody above him, so the pressure on his body is 

very great. At each higher level, the overlying weight is less, 

so the pressure decreases. Gravitational equilibrium in the 

Sun means that the pressure increases with depth, making 

the Sun extremely hot and dense in its central core (Figure 

15.2). 

THINK ABOUT IT 

Earth’s atmosphere is also in gravitational equilibrium, with the 

weight of upper layers supported by the pressure in lower 

layers. Use this idea to explain why the air gets thinner at higher 

altitudes. 

chapter 15 • Our Star 497

pressure 

gravity 

Figure 15.2 Gravitational equilibrium in the Sun: At each point 

inside, the pressure pushing outward balances the weight of the 

overlying layers. 

Figure 15.1 An acrobat stack is in gravitational equilibrium: The 

lowest person supports the most weight and feels the greatest 

pressure, and the overlying weight and underlying pressure decrease 

for those higher up. 

Although the Sun today maintains its gravitational 

equilibrium with energy generated by nuclear fusion, the 

energy-generation mechanism of gravitational contraction 

was important in the distant past and will be important 

again in the distant future. Our Sun was born from a collapsing 

cloud of interstellar gas. The contraction of the 

cloud released gravitational potential energy, raising the 

interior temperature higher and higher—but not high 

enough to halt the contraction. The cloud continued to 

shrink because thermal radiation from the cloud’s surface 

carried away much of the energy released by contraction, 

even while the interior temperature was rising. When the 

central temperature and density eventually reached the 

values necessary to sustain nuclear fusion, energy generation 

in the Sun’s interior matched the energy lost from the 

surface in the form of radiation. With the onset of fusion, 

the Sun entered a long-lasting state of gravitational equilibrium 

that has persisted for the last 4.6 billion years. 

About 5 billion years from now, when the Sun finally 

exhausts its nuclear fuel, the internal pressure will drop, and 

gravitational contraction will begin once again. As we will see 

later, some of the most important and spectacular processes 

in astronomy hinge on this ongoing “battle” between the 

crush of gravity and a star’s internal sources of pressure. 

In summary, the answer to the question “Why does the 

Sun shine?” is that about 4.6 billion years ago gravitational 

contraction made the Sun hot enough to sustain nuclear 

fusion in its core. Ever since, energy liberated by fusion has 

maintained the Sun’s gravitational equilibrium and kept the 

Sun shining steadily, supplying the light and heat that sustain 

life on Earth. 

15.2 Plunging to the 

Center of the Sun: 


In the rest of this chapter, we will discuss in detail how 

the Sun produces energy and how that energy travels to 

Earth. First, to get a “big picture” view of the Sun, let’s 

imagine you have a spaceship that can somehow withstand 

the immense heat and pressure of the solar interior 

and take an imaginary journey from Earth to the center 

of the Sun. 

Approaching the Surface 

As you begin your voyage from Earth, the Sun appears as a 

whitish ball of glowing gas. With spectroscopy [Section 7.3], 

you verify that the Sun’s mass is 70% hydrogen and 28% 

helium. Heavier elements make up the remaining 2%. 

The total power output of the Sun, called its luminosity, 

is an incredible 3.8 10 26 watts. That is, every second, 

the Sun radiates a total of 3.8 10 26 joules of energy into 

space (recall that 1 watt 1 joule/s). If we could somehow 

capture and store just 1 second’s worth of the Sun’s lumi- 

498 part V • Stellar Alchemy

Table 15.1 Basic Properties of the Sun 

Radius (R Sun 

) 

Mass (M Sun 

) 

Luminosity (L Sun 

) 

Composition (by 

percentage of mass) 

Rotation rate 

Surface temperature 

Core temperature 

696,000 km (about 109 times the 

radius of Earth) 

2 10 30 kg (about 300,000 times 

the mass of Earth) 

3.8 10 26 watts 

70% hydrogen, 28% helium, 

2% heavier elements 

25 days (equator) to 30 days (poles) 

5,800 K (average); 4,000 K (sunspots) 

15 million K 

Figure 15.3 This photo of the 

visible surface of the Sun shows 

several dark sunspots. 

nosity, it would be enough to meet current human energy 

demands for roughly the next 500,000 years! 

Of course, only a tiny fraction of the Sun’s total energy 

output reaches Earth, with the rest dispersing in all directions 

into space. Most of this energy is radiated in the form 

of visible light, but once you leave the protective blanket of 

Earth’s atmosphere you’ll encounter significant amounts of 

other types of solar radiation, including dangerous ultraviolet 

and X rays. Your spaceship will require substantial 

shielding to protect you from serious radiation burns caused 

by these high-energy forms of light. 

Through a telescope, you can see that the Sun seethes 

with churning gases. At most times you’ll detect at least 

a few sunspots blotching its surface (Figure 15.3). If you 

focus your telescope solely on a sunspot, you’ll find that it is 

blindingly bright. Sunspots appear dark only in contrast to 

the even brighter solar surface that surrounds them. A typical 

sunspot is large enough to swallow the entire Earth, dramatically 

illustrating that the Sun is immense by any 

earthly standard. The Sun’s radius is nearly 700,000 kilometers, 

and its mass is 2 10 30 kilograms—about 300,000 

times more massive than Earth. 

Sunspots appear to move from day to day along with 

the Sun’s rotation. If you watch very carefully, you may 

notice that sunspots near the solar equator circle the Sun 

faster than those at higher solar latitudes. This observation 

reveals that, unlike a spinning ball, the entire Sun does 

not rotate at the same rate. Instead, the solar equator completes 

one rotation in about 25 days, and the rotation period 

increases with latitude to about 30 days near the solar 

poles. Table 15.1 summarizes some of the basic properties 



VIS 

As a brief review, describe how we measure the mass of the 

Sun using Newton’s version of Kepler’s third law. (Hint: Look 

back at Chapter 5.) 

As you and your spaceship continue to fall toward the 

Sun, you notice an increasingly powerful headwind exerting 

a bit of drag on your descent. This headwind, called the 

solar wind, is created by ions and subatomic particles flowing 

outward from the solar surface. The solar wind helps 

shape the magnetospheres of planets [Sections 11.3, 12.4] 

and blows back the material that forms the tails of comets 

[Section 13.4]. 

A few million kilometers above the solar surface, you 

enter the solar corona, the tenuous uppermost layer of the 

Sun’s atmosphere (Figure 15.4). Here you find the temperature 

to be astonishingly high—about 1 million Kelvin. This 

region emits most of the Sun’s X rays. However, the density 

here is so low that your spaceship feels relatively little heat 

despite the million-degree temperature [Section 4.2]. 

Nearer the surface, the temperature suddenly drops to 

about 10,000 K in the chromosphere, the primary source 

of the Sun’s ultraviolet radiation. At last you plunge through 

the visible surface of the Sun, called the photosphere,where 

the temperature averages just under 6,000 K. Although the 

photosphere looks like a well-defined surface from Earth, 

it consists of gas far less dense than Earth’s atmosphere. 

Throughout the solar atmosphere, you notice that 

the Sun has its own version of weather, in which conditions 

at a particular altitude differ from one region to another. 

Some regions of the chromosphere and corona are particularly 

hot and bright, while other regions are cooler and 

less dense. In the photosphere, sunspots are cooler than 

the surrounding surface, though they are still quite hot 

and bright by earthly standards. In addition, your compass 

goes crazy as you descend through the solar atmosphere, 

indicating that solar weather is shaped by intense magnetic 

fields. Occasionally, huge magnetic storms occur, shooting 

hot gases far into space. 

Into the Sun 

Up to this point in your journey, you may have seen Earth 

and the stars when you looked back, but as you slip beneath 

the photosphere, blazing light engulfs you. You are 


solar wind 

photosphere 

corona 

chromosphere 

convection 

zone 

core 

radiation 

zone 

solar wind 

Figure 15.4 The basic structure of the Sun. Nuclear fusion in the solar core generates the Sun’s energy. 

Photons of light carry that energy through the radiation zone to the bottom of the convection zone. Rising 

plumes of hot gas then transport the energy through the convection zone to the photosphere, where it is 

radiated into space. The photosphere, at a temperature of roughly 6,000 K, is relatively cool compared to 

the layers that lie above it. The temperature of the chromosphere, which is directly above the photosphere, 

exceeds 10,000 K. The temperature of the corona, extending outward from the chromosphere, can reach 

1 million degrees. Because the coronal gas is so hot, some of it escapes the Sun’s gravity, forming a solar 

wind that blows past Earth and out beyond Pluto. 

inside the Sun, and your spacecraft is tossed about by incredible 

turbulence. If you can hold steady long enough to 

see what is going on around you, you’ll notice spouts of 

hot gas rising upward, surrounded by cooler gas cascading 

down from above. You are in the convection zone,where 

energy generated in the solar core travels upward, transported 

by the rising of hot gas and falling of cool gas called 

convection [Section 10.2].With some quick thinking, you 

may realize that the photosphere above you is the top of 

the convection zone and that convection is the cause of 

the Sun’s seething, churning appearance. 

As you descend through the convection zone, the surrounding 

density and pressure increase substantially, along 

with the temperature. Soon you reach depths at which the 

Sun is far denser than water. Nevertheless, it is still a gas 

(more specifically, a plasma of positively charged ions and 

free electrons) because each particle moves independently 

of its neighbors [Section 4.3]. 

About a third of the way down to the center, the turbulence 

of the convection zone gives way to the calmer 

plasma of the radiation zone,where energy is carried outward 

primarily by photons of light. The temperature rises 

to almost 10 million K, and your spacecraft is bathed in 

X rays trillions of times more intense than the visible light 

at the solar surface. 


COMMON MISCONCEPTIONS 

The Sun Is Not on Fire 

fission 

fusion 

We are accustomed to saying that the Sun is “burning,” 

a way of speaking that conjures up images of a giant bonfire 

in the sky. However, the Sun does not burn in the 

same sense as a fire burns on Earth. Fires on Earth generate 

light through chemical changes that consume oxygen 

and produce a flame. The glow of the Sun has more 

in common with the glowing embers left over after the 

flames have burned out. Much like hot embers, the Sun’s 

surface shines with the visible thermal radiation produced 

by any object that is sufficiently hot [Section 6.4]. 

However, hot embers quickly stop glowing as they 

cool, while the Sun keeps shining because its surface is 

kept hot by the energy rising from the Sun’s core. Because 

this energy is generated by nuclear fusion, we 

sometimes say that it is the result of “nuclear burning”— 

a term that suggests nuclear changes in much the same 

way that “chemical burning” suggests chemical changes. 

Nevertheless, while it is reasonable to say that the Sun 

undergoes nuclear burning in its core, it is not accurate 

to speak of any kind of burning on the Sun’s surface, 

where light is produced primarily by thermal radiation. 

No real spacecraft could survive, but your imaginary 

one keeps plunging straight down to the solar core.There 

you finally find the source of the Sun’s energy: nuclear fusion 

transforming hydrogen into helium. At the Sun’s center, 

the temperature is about 15 million K, the density is more 

than 100 times that of water, and the pressure is 200 billion 

times that on the surface of Earth. The energy produced 

in the core today will take about a million years to 

reach the surface. 

With your journey complete, it’s time to turn around 

and head back home. We’ll continue this chapter by studying 

fusion in the solar core and then tracing the flow of the 

energy generated by fusion as it moves outward through 

the Sun. 


The prospect of turning common metals like lead into 

gold enthralled those who pursued the medieval practice 

of alchemy. Sometimes they tried primitive scientific approaches, 

such as melting various ores together in a vessel 

called a crucible. Other times they tried magic. Their getrich-quick 

schemes never managed to work. Today we 

know that there is no easy way to turn other elements into 

gold, but it is possible to transmute one element or isotope 

into another. 

If a nucleus gains or loses protons, its atomic number 

changes and it becomes a different element. If it gains or 

Figure 15.5 Nuclear fission splits a nucleus into smaller nuclei 

(not usually of equal size), while nuclear fusion combines smaller 

nuclei into a larger nucleus. 

loses neutrons, its atomic mass changes and it becomes a 

different isotope [Section 4.3].The process of splitting a nucleus 

into two smaller nuclei is called nuclear fission.The 

process of combining nuclei to make a nucleus with a greater 

number of protons or neutrons is called nuclear fusion 

(Figure 15.5). Human-built nuclear power plants rely 

on nuclear fission of uranium or plutonium. The nuclear 

power plant at the center of the Sun relies on nuclear fusion, 

turning hydrogen into helium. 

Nuclear Fusion 

The 15 million K plasma in the solar core is like a “soup” 

of hot gas, with bare, positively charged atomic nuclei (and 

negatively charged electrons) whizzing about at extremely 

high speeds. At any one time, some of these nuclei are on 

high-speed collision courses with each other. In most cases, 

electromagnetic forces deflect the nuclei, preventing actual 

collisions, because positive charges repel one another. If 

nuclei collide with sufficient energy, however, they can stick 

together to form a heavier nucleus (Figure 15.6). 

Sticking positively charged nuclei together is not easy. 

The strong force,which binds protons and neutrons together 

in atomic nuclei, is the only force in nature that can 

At low speeds, electromagnetic 

repulsion prevents the collision 

of nuclei. 

At high speeds, nuclei come close 

enough for the strong force to bind 

them together. 

Figure 15.6 Positively charged nuclei can fuse only if a highspeed 

collision brings them close enough for the strong force 

to come into play. 


p 

p 

Step 1 Step 2 Step 3 

p 

n 

e 

gamma ray 

Key: 

e 

electron 

gamma 

ray 

e e 

positron p proton 

Total reaction 

p 

p 

n 

p 

p 

p 

p 

n p 

gamma ray 

n p 

p 

p 

p 

n n 

p 

n 

p 

p 

p 

e 

n p 

p 

e 

gamma ray 

gamma ray 

neutrino 

gamma ray 

p n 

n 

neutron 

Figure 15.7 Hydrogen fuses into helium in the Sun by way of the proton–proton chain. In step 1, two 

protons fuse to create a deuterium nucleus consisting of a proton and a neutron. In step 2, the deuterium 

nucleus and a proton fuse to form helium-3, a rare form of helium. In step 3, two helium-3 nuclei fuse 

to form helium-4, the common form of helium. 

overcome the electromagnetic repulsion between two positively 

charged nuclei [Section S4.2].In contrast to gravitational 

and electromagnetic forces, which drop off gradually 

as the distances between particles increase (by an inverse 

square law [Section 5.3]), the strong force is more like glue 

or Velcro: It overpowers the electromagnetic force over 

very small distances but is insignificant when the distances 

between particles exceed the typical sizes of atomic nuclei. 

The trick to nuclear fusion, therefore, is to push the positively 

charged nuclei close enough together for the strong 

force to outmuscle electromagnetic repulsion. 

The high pressures and temperatures in the solar core 

are just right for fusion of hydrogen nuclei into helium 

nuclei. The high temperature is important because the nuclei 

must collide at very high speeds if they are to come 

close enough together to fuse. (Quantum tunneling is also 

important to this process [Section S4.5].) The higher the 

temperature, the harder the collisions, making fusion reactions 

more likely at higher temperatures. The high pressure 

of the overlying layers is necessary because without it the 

hot plasma of the solar core would simply explode into space, 

shutting off the nuclear reactions. In the Sun, the pressure 

is high and steady, allowing some 600 million tons of hydrogen 

to fuse into helium every second. 

Hydrogen Fusion in the Sun: 

The Proton–Proton Chain 

Recall that hydrogen nuclei are nothing more than individual 

protons, while the most common form of helium 

consists of two protons and two neutrons. Thus, the 

overall hydrogen fusion reaction in the Sun is: 

p 

p 

4 1 H 

p 

p 

p n n 

p 

1 4 He 

energy 

However, collisions between two nuclei are far more 

common than three- or four-way collisions, so this overall 

reaction proceeds through steps that involve just two nuclei 

at a time. The sequence of steps that occurs in the Sun is 

called the proton–proton chain because it begins with collisions 

between individual protons (hydrogen nuclei). 

Figure 15.7 illustrates the steps in the proton–proton chain: 

Step 1. Two protons fuse to form a nucleus consisting of 

one proton and one neutron, which is the isotope of hydrogen 

known as deuterium. Note that this step converts a 

proton into a neutron, reducing the total nuclear charge 

from 2 for the two fusing protons to 1 for the resulting 

deuterium nucleus. The lost positive charge is carried off 

by a positron, the antimatter version of an electron with 

a positive rather than negative charge [Section S4.2].A neutrino—a 

subatomic particle with a very tiny mass—is also 

produced in this step.* The positron won’t last long, because 

it soon meets up with an ordinary electron, resulting 

*Producing a neutrino is necessary because of a law called conservation 

of lepton number: The number of leptons (e.g., electrons or neutrinos 

[Chapter S4]) must be the same before and after the reaction. The lepton 

number is zero before the reaction because there are no leptons. Among 

the reaction products, the positron (antielectron) has lepton number 1 

because it is antimatter, and the neutrino has lepton number 1. Thus, 

the total lepton number remains zero. 


in the creation of two gamma-ray photons through matter– 

antimatter annihilation. 

Step 2. A fair number of deuterium nuclei are always present 

along with the protons and other nuclei in the solar 

core, since step 1 occurs so frequently in the Sun (about 

10 38 times per second). Step 2 occurs when one of these 

deuterium nuclei collides and fuses with a proton. The 

result is a nucleus of helium-3, a rare form of helium with 

two protons and one neutron. This reaction also produces 

a gamma-ray photon. 

Step 3. The third and final step of the proton–proton chain 

requires the addition of another neutron to the helium-3, 

thereby making normal helium-4. This final step can proceed 

in several different ways, but the most common route 

involves a collision of two helium-3 nuclei. Each of these 

helium-3 nuclei resulted from a prior, separate occurrence 

of step 2 somewhere in the solar core. The final result is a 

normal helium-4 nucleus and two protons. 

Total reaction. Somewhere in the solar core, steps 1 and 2 

must each occur twice to make step 3 possible. Six protons 

go into each complete cycle of the proton–proton chain, but 

two come back out. Thus, the overall proton–proton chain 

converts four protons (hydrogen nuclei) into a helium-4 

nucleus, two positrons, two neutrinos, and two gamma rays. 

Each resulting helium-4 nucleus has a mass that is 

slightly less (by about 0.7%) than the combined mass of 

the four protons that created it. Overall, fusion in the Sun 

converts about 600 million tons of hydrogen into 596 million 

tons of helium every second. The “missing” 4 million 

tons of matter becomes energy in accord with Einstein’s 

formula E mc 2 .About 98% of the energy emerges as 

kinetic energy of the resulting helium nuclei and radiative 

energy of the gamma rays. As we will see, this energy 

slowly percolates to the solar surface, eventually emerging 

as the sunlight that bathes Earth. About 2% of the energy 

is carried off by the neutrinos. Neutrinos rarely interact 

with matter (because they respond only to the weak force 

[Section S4.2]), so most of the neutrinos created by the 

proton–proton chain pass straight from the solar core 

through the solar surface and out into space. 

The Solar Thermostat 

The rate of nuclear fusion in the solar core, which determines 

the energy output of the Sun, is very sensitive to 

temperature. A slight increase in temperature would mean 

a much higher fusion rate, and a slight decrease in temperature 

would mean a much lower fusion rate. If the 

Sun’s rate of fusion varied erratically, the effects on Earth 

might be devastating. Fortunately, the Sun’s central temperature 

is steady thanks to gravitational equilibrium—the 

balance between the pull of gravity and the push of internal 

pressure. 

Outside the solar core, the energy produced by fusion 

travels toward the Sun’s surface at a slow but steady rate. In 

this steady state, the amount of energy leaving the top of 

each gas layer within the Sun precisely balances the energy 

entering from the bottom (Figure 15.8). Suppose the core 

temperature of the Sun rose very slightly. The rate of nuclear 

fusion would soar, generating lots of extra energy. Because 

energy moves so slowly through the Sun, this extra 

energy would be bottled up in the core, causing an increase 

in the core pressure. The push of this pressure would temporarily 

exceed the pull of gravity, causing the core to expand 

and cool. With cooling, the fusion rate would drop 

back down. The expansion and cooling would continue until 

gravitational equilibrium was restored, at which point the 

fusion rate would return to its original value. 

An opposite process would restore the normal fusion 

rate if the core temperature dropped. A decrease in core 

temperature would lead to decreased nuclear burning, 

a drop in the central pressure, and contraction of the core. 

As the core shrank, its temperature would rise until the 

burning rate returned to normal. 

The response of the core pressure to changes in the 

nuclear fusion rate is essentially a thermostat that keeps 

the Sun’s central temperature steady. Any change in 

the core temperature is automatically corrected by the 

change in the fusion rate and the accompanying change 

in pressure. 

While the processes involved in gravitational equilibrium 

prevent erratic changes in the fusion rate, they also 

ensure that the fusion rate gradually rises over billions of 

years. Because each fusion reaction converts four hydrogen 

nuclei into one helium nucleus, the total number of independent 

particles in the solar core is gradually falling. This 

gradual reduction in the number of particles causes the 

solar core to shrink. 

The slow shrinking of the solar core means that it must 

generate energy more rapidly to counteract the stronger 

compression of gravity, so the solar core gradually gets 

hotter as it shrinks. Theoretical models indicate that the 

Sun’s core temperature should have increased enough to 

raise its fusion rate and the solar luminosity by about 30% 

since the Sun was born 4.6 billion years ago. 

How did the gradual increase in solar luminosity affect 

Earth? Geological evidence shows that Earth’s surface temperature 

has remained fairly steady since Earth finished 

forming more than 4 billion years ago, despite this 30% 

increase in the Sun’s energy output, because Earth has its 

own thermostat. This “Earth thermostat” is the carbon 

dioxide cycle. By maintaining a fairly steady level of atmospheric 

carbon dioxide, the carbon dioxide cycle regulates 

the greenhouse effect that maintains Earth’s surface temperature 

[Section 14.4]. 

“Observing” the Solar Interior 

We cannot see inside the Sun, so you may be wondering 

how we can know so much about what goes on underneath 

its surface. Astronomers can study the Sun’s interior in 

three different ways: through mathematical models of the 


large decrease in 

rate of fusion 

slight decrease in 

core temperature 

slight rise in 

core temperature 

large rise in 

rate of fusion 

Solar 

Thermostat: 

Gravitational 

Equilibrium 

Because the energy 

supply is diminished, 

gravity starts to 

overcome thermal 

pressure. 

Increased energy 

output enables 

thermal pressure 

to overcome 

gravity. 

Gravity compresses 

the core, heats it up, 

and restores fusion 

rate to normal value. 

Increased thermal 

pressure causes the 

core to expand and then 

cool, which restores fusion 

rate to normal value. 

Figure 15.8 The solar thermostat. Gravitational equilibrium regulates the Sun’s core temperature. Everything 

is in balance if the amount of energy leaving the core equals the amount of energy produced by 

fusion. A rise in core temperature triggers a chain of events that causes the core to expand, lowering its 

temperature to its normal value. A decrease in core temperature triggers the opposite chain of events, 

also restoring the normal core temperature. 

Sun, observations of “sun quakes,” and observations of 

solar neutrinos. 

Mathematical Models The primary way we learn about 

the interior of the Sun and other stars is by creating mathematical 

models that use the laws of physics to predict the 

internal conditions. A basic model uses the Sun’s observed 

composition and mass as inputs to equations that describe 

gravitational equilibrium, the solar thermostat, and the 

rate at which solar energy moves from the core to the photosphere. 

With the aid of a computer, we can use the model 

to calculate the Sun’s temperature, pressure, and density 

at any depth. We can then predict the rate of nuclear fusion 

in the solar core by combining these calculations with 

knowledge about nuclear fusion gathered in laboratories 

here on Earth. 

Remarkably, such models correctly “predict” the radius, 

surface temperature, luminosity, age, and many other 

properties of the Sun. However, current models do not 

predict everything about the Sun correctly. Scientists are 

constantly working to discover what is missing from them. 

Successful prediction of so many observed characteristics 

of the Sun gives us confidence that the models are on the 

right track and that we really do understand what is going 

on inside the Sun. 

Sun Quakes A second way to learn about the inside of 

the Sun is to observe “sun quakes”—vibrations of the Sun 

that are similar to the vibrations of the Earth caused by 

earthquakes, although they are generated very differently. 

Earthquakes occur when Earth’s crust suddenly shifts, generating 

seismic waves that propagate through Earth’s interior 

[Section 10.2].We can learn about Earth’s interior by recording 

seismic waves on Earth’s surface with seismographs. 

Sun quakes result from waves of pressure (sound waves) 

that propagate deep within the Sun at all times. These 

waves cause the solar surface to vibrate when they reach 

it. Although we cannot set up seismographs on the Sun, 

we can detect the vibrations of the surface by measuring 

Doppler shifts [Section 6.5]. Light from portions of the surface 

that are rising toward us is slightly blueshifted, while 

light from portions that are falling away from us is slightly 

redshifted. The vibrations are relatively small but measurable 

(Figure 15.9). 


Figure 15.9 Vibrations on the surface of 

the Sun can be detected by Doppler shifts. 

In this schematic representation, red indicates 

falling gas, and blue indicates rising gas. 

The speckled region indicates the convection 

zone. The vibration pattern illustrated 

here is just one of many possible patterns. 

The overall vibration pattern of the Sun is a 

complex combination of patterns similar to 

this one. 

In principle, we can deduce a great deal about the 

solar interior by carefully analyzing these vibrations. (By 

analogy to seismology on Earth, this type of study of the 

Sun is called helioseismology—helios means “sun.”) Results 

to date confirm that our mathematical models of the 

solar interior are on the right track (Figure 15.10). At the 

same time, they provide data that can be used to improve 

the models further. 

Mathematical Insight 15.1 Mass-Energy Conversion in the Sun 

We can calculate how much mass the Sun loses through nuclear 

fusion by comparing the input and output masses of the proton– 

proton chain. A single proton has a mass of 1.6726 10 27 kg, 

so four protons have a mass of 6.690 10 27 kg. 

A helium-4 nucleus has a mass of only 6.643 10 27 kg, 

slightly less than the mass of the four protons. The difference is: 

6.690 10 27 kg 6.643 10 27 kg 4.7 10 29 kg 

which is 0.7%, or 0.007, of the original mass. Thus, for example, 

when 1 kilogram of hydrogen fuses, the resulting helium weighs 

only 993 grams, while 7 grams of mass turns into energy. 

To calculate the total amount of mass converted to energy 

in the Sun each second, we use Einstein’s equation E mc 2 .The 

total energy produced by the Sun each second is 3.8 10 26 joules, 

so we can solve for the total mass converted to energy each second: 

E mc 2 E 

⇒ m c 2 

 

3.8 10 26 joules 

 

 

3.0 10 8 m s 2 

4.2 10 9 kg 

The Sun loses about 4 billion kilograms of mass every second, 

which is roughly equivalent to the combined mass of nearly 

100 million people. 

Example: How much hydrogen is converted to helium each 

second in the Sun? 

Solution: We have already calculated that the Sun loses 4.2 10 9 kg 

of mass each second and that this is only 0.7% of the mass of hydrogen 

that is fused: 

4.2 10 9 kg 0.007 mass of hydrogen fused 

We now solve for the mass of hydrogen fused: 

mass of 

hydrogen fused 4.2 10 

9 kg 

 

0.007 

6.0 10 11 kg 1m etric 

ton 

103 

kg 

6.0 10 8 metric tons 

The Sun fuses 600 million metric tons of hydrogen each second, 

of which about 4 million tons becomes energy. The remaining 

596 million tons becomes helium. 


Key 

model 

10 7 data 

Key 

100 model 

data 

10 

temperature (K) 

10 6 

core 

radiation 

zone 

convection 

zone 

density (g/cm 3 ) 

1 

0.1 

core 

radiation 

zone 

convection 

zone 

0.01 

0 0.2 0.4 0.6 0.8 1.0 

fraction of the Sun’s radius 

a Temperature at different radii within the Sun. 

0 0.2 0.4 0.6 0.8 1.0 

fraction of the Sun’s radius 

b Density at different radii within the Sun. (The density of water is 

1 g/cm 3 .) 

Figure 15.10 Agreement between mathematical models of solar structure and actual measurements of 

solar structure derived from “sun quakes.” The red lines show predictions of mathematical models of the 

Sun. The blue lines show the interior structure of the Sun as indicated by vibrations of the Sun’s surface. 

These vibrations tell us about conditions deep within the Sun because they are produced by sound waves 

that propagate through the Sun’s interior layers. 

Solar Neutrinos Another way to study the solar interior 

is to observe the neutrinos coming from fusion reactions in 

the core. Don’t panic, but as you read this sentence about a 

thousand trillion solar neutrinos will zip through your body. 

Fortunately, they won’t do any damage, because neutrinos 

rarely interact with anything. Neutrinos created by fusion 

in the solar core fly quickly through the Sun as if passing 

through empty space. In fact, while an inch of lead will stop 

an X ray, stopping an average neutrino would require a slab 

of lead more than 1 light-year thick! Clearly, counting neutrinos 

is dauntingly difficult, because virtually all of them 

stream right through any detector built to capture them. 


Is the number of solar neutrinos zipping through our bodies 

significantly lower at night? (Hint: How does the thickness of 

Earth compare with the thickness of a slab of lead needed 

to stop an average neutrino?) 

Nevertheless, neutrinos do occasionally interact with 

matter, and it is possible to capture a few solar neutrinos 

with a large enough detector. Neutrino detectors are usually 

placed deep inside mines so that the overlying layers of 

rock block all other kinds of particles coming from outer 

space except neutrinos, which pass through rock easily. The 

first major solar neutrino detector, built in the 1960s, was 

located 1,500 meters underground in the Homestake gold 

mine in South Dakota (Figure 15.11). 

The detector for this “Homestake experiment” consisted 

of a 400,000-liter vat of chlorine-containing dry-cleaning 

fluid. It turns out that, on very rare occasions, a chlorine 

nucleus can capture a neutrino and change into a nucleus 

of radioactive argon. By looking for radioactive argon in 

the tank of cleaning fluid, experimenters could count the 

number of neutrinos captured in the detector. 

From the many trillions of solar neutrinos that passed 

through the tank of cleaning fluid each second, experimenters 

expected to capture an average of just one neutrino 

per day. This predicted capture rate was based on measured 

properties of chlorine nuclei and models of nuclear fusion 

in the Sun. However, over a period of more than two 

decades, neutrinos were captured only about once every 

3 days on average. That is, the Homestake experiment detected 

only about one-third of the predicted number of 

neutrinos. This disagreement between model predictions 

and actual observations came to be called the solar neutrino 

problem. 

The shortfall of neutrinos found with the Homestake 

experiment led to many more recent attempts to detect solar 

neutrinos using more sophisticated detectors (Figure 15.12). 

The chlorine nuclei in the Homestake experiment could 


Figure 15.11 This tank of dry-cleaning fluid (visible underneath 

the catwalk), located deep within South Dakota’s Homestake mine, 

was a solar neutrino detector. The chlorine nuclei in the cleaning 

fluid turned into argon nuclei when they captured neutrinos from 

the Sun. 

a Scientists inspecting individual detectors within Super-Kamiokande. 

capture only high-energy neutrinos that are produced by 

one of the rare pathways of step 3 in the proton–proton 

chain (not shown in Figure 15.7). More recent experiments 

can detect lower-energy neutrinos, including those produced 

by step 1 of the proton–proton chain, and therefore offer a 

better probe of fusion in the Sun. To date, all these experiments 

have found fewer neutrinos than current models of 

the Sun predict. This discrepancy between model and experiment 

probably means one of two things: Either something 

is wrong with our models of the Sun, or something is 

missing in our understanding of how neutrinos behave. 


Although the observed number of neutrinos falls short of 

theoretical predictions, experiments like Homestake have 

shown that at least some neutrinos are coming from the 

Sun. Explain why this provides direct evidence that nuclear 

fusion really is taking place in the Sun right now. (Hint: See 

Figure 15.7.) 

For the moment, many physicists and astronomers 

are betting that we understand the Sun just fine and that 

the discrepancy has to do with the neutrinos themselves. 

One intriguing idea arises from the fact that neutrinos come 

in three types: electron neutrinos, muon neutrinos, and 

tau neutrinos [Section S4.2]. 

Fusion reactions in the Sun produce only electron neutrinos, 

and most solar neutrino detectors can detect only 

b An image of the Sun created by tracing the paths of 

neutrinos detected by Super-Kamiokande back to the Sun. 

Figure 15.12 The Super-Kamiokande experiment in Japan is one 

of the world’s premier neutrino detectors. 


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. 


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 


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. 


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 


a This close-up view of the Sun’s surface 

(right) shows two large sunspots and several 

smaller ones. Both of the big sunspots are 

roughly as large as Earth. 

Figure 15.15 Sunspots are regions of intense 

magnetic activity. 

VIS 

tight magnetic field lines suppress convection within each 

sunspot and prevent surrounding plasma from sliding sideways 

into the sunspot. With hot plasma unable to enter the 

region, the sunspot plasma becomes cooler than that of the 

rest of the photosphere (Figure 15.17a). 

The magnetic field lines connecting two sunspots often 

soar high above the photosphere, through the chromosphere, 

and into the corona (Figure 15.17b). These vaulted loops 

of magnetic field sometimes appear as solar prominences, 

in which the field traps gas that may glow for days or even 

weeks. Some prominences rise to heights of more than 

100,000 kilometers above the Sun’s surface (Figure 15.18). 

The most dramatic events on the solar surface are 

solar flares,which emit bursts of X rays and fast-moving 

charged particles into space (Figure 15.19). Flares generally 

occur in the vicinity of sunspots, leading us to suspect that 

they occur when the magnetic field lines become so twisted 

and knotted that they can no longer bear the tension. The 

magnetic field lines suddenly snap like tangled elastic bands 

twisted beyond their limits, releasing a huge amount of 

energy. This energy heats the nearby plasma to 100 million 

K over the next few minutes to few hours, generating X rays 

and accelerating some of the charged particles to nearly the 

speed of light. 

The Chromosphere and Corona 

The high temperatures of the chromosphere and corona 

perplexed scientists for decades. After all, temperatures 

gradually decline as we move outward from the core to the 

VIS 

b Spectra of sunspots can be 

used to measure the strength 

of their magnetic fields. This 

image shows the spectrum 

of a sunspot and its surroundings. 

The sunspot region 

shows up as dark horizontal 

bands because it is darker 

than the rest of the solar surface 

in its vicinity. The vertical 

bands are absorption lines 

that are present both inside 

and outside the sunspots. 

The influence of strong magnetic 

fields within the sunspot 

region splits a single absorption 

line into three parts. 

Measuring the separation 

between these lines tells us 

the strength of the magnetic 

field within the sunspot. 

top of the photosphere. Why should this decline suddenly 

reverse? Some aspects of this atmospheric heating remain 

a mystery today, but we have at least a general explanation: 

The Sun’s strong magnetic fields carry energy upward from 

the churning solar surface to the chromosphere and corona. 

More specifically, the rising and falling of gas in the 

convection zone probably shakes magnetic field lines beneath 

the solar surface. This shaking generates waves along the 

magnetic field lines that carry energy upward to the solar 


weaker magnetic field 

stronger 

magnetic 

field 

e 

e 

e 

weaker magnetic field 

a Magnetic field lines follow b Lines closer together c Charged particles follow 

the directions that compass indicate a stronger field. paths that spiral along 

needles would point. 

magnetic field lines. 

e 

Figure 15.16 We draw magnetic field 

lines to represent invisible magnetic fields. 

atmosphere. Precisely how the waves deposit their energy 

in the chromosphere and corona is not known, but the 

waves agitate the low-density plasma of these layers, somehow 

heating them to high temperatures. Much of this heating 

appears to happen near where the magnetic field lines 

emerge from the Sun’s surface. 

According to this model of solar heating, the same 

magnetic fields that keep sunspots cool make the overlying 

plasma of the chromosphere and corona hot. We can test 

this idea observationally. The gas of the chromosphere and 

corona is so tenuous that we cannot see it with our eyes 

except during a total eclipse, when we can see the faint visible 

light scattered by electrons in the corona [Section 2.5]. 

However, the roughly 10,000 K plasma of the chromosphere 

emits strongly in the ultraviolet, and the million K plasma 

of the corona is the source of virtually all X rays coming 

Magnetic fields 

trap gas. 

T 5,800 K 

sunspots 

T 4,500 K 

T 5,800 K 

convection 

cells 

Magnetic fields of sunspots suppress convection 

and prevent surrounding plasma from sliding 

sideways into sunspot. 

a Pairs of sunspots are connected by tightly wound magnetic 

field lines. 

Figure 15.17 Loops of magnetic field lines can arch high above 

the solar surface, reaching heights many times larger than Earth’s 

diameter. 

b This X-ray photo (from NASA’s 

TRACE mission) shows gas trapped 

within looped magnetic field lines. 

X-ray 


Figure 15.18 A gigantic solar prominence 

erupts from the solar surface at the upper 

right of this ultraviolet-light photo (from the 

SOHO mission). The gas within this prominence, 

which is over 20 times the size of 

Earth, is quite hot but still cooler than the 

million-degree gas of the surrounding corona. 

UV 

Figure 15.19 This photo (from 

TRACE) of ultraviolet light emitted 

by hydrogen atoms shows a solar 

flare erupting from the Sun’s surface. 

UV 

from the Sun. Figure 15.20 shows an X-ray image of the Sun. 

As the solar heating model predicts, the brightest regions 

of the corona tend to be directly above sunspot groups. 

Some regions of the corona, called coronal holes, barely 

show up in X-ray images. More detailed analyses show that 

the magnetic field lines in coronal holes project out into 

space like broken rubber bands, allowing particles spiraling 

along them to escape the Sun altogether. These particles 

streaming outward from the corona constitute the solar 

wind, which blows through the solar system at an average 

speed of about 500 kilometers per second and has important 

effects on planetary surfaces, atmospheres, and magnetospheres. 

Well beyond the planets, the pressure of interstellar 

gas must eventually halt the solar wind. The Pioneer 

and Voyager spacecraft that visited the outer planets in the 

1970s and 1980s are still traveling outward from our solar 

system and may soon encounter this “boundary” (called the 

heliopause) of the realm of the Sun. 

The solar wind also gives us something tangible to 

study. In the same way that meteorites provide us with samples 

of asteroids we’ve never visited, solar wind particles 

captured by satellites provide us with a sample of material 

from the Sun. Analysis of these solar particles has reassuringly 

verified that the Sun is made mostly of hydrogen, 

just as we conclude from studying the Sun’s spectrum. 


Figure 15.20 An X-ray image of the 

Sun reveals the million-degree gas of the 

corona. Brighter regions of this image 

(yellow) correspond to regions of stronger 

X-ray emission. The darker regions are 

the coronal holes from which the solar 

wind escapes. (From the Yohkoh space 

observatory.) 

X-ray 

15.5 Solar Weather 

and Climate 

Individual sunspots, prominences, and flares are short-lived 

phenomena, somewhat like storms on Earth. They constitute 

what we call solar weather or solar activity.You know 

from personal experience that the Earth’s weather is notoriously 

unpredictable. The same is true for the Sun: We 

cannot predict precisely when or where a particular sunspot 

or flare will appear. Earth’s climate, on the other hand, is 

quite regular from season to season. So it is with the Sun, 

where despite day-to-day variations the general nature and 

intensity of solar activity follow a predictable cycle. 

The Sunspot Cycle 

Long before we realized that sunspots were magnetic disturbances, 

astronomers had recognized patterns in sunspot 

activity. The most notable pattern is the number of sunspots 

visible on the Sun at any particular time. Thanks to 

telescopic observations of the Sun recorded by astronomers 

since the 1600s, we know that the number of sunspots 

gradually rises and falls in a sunspot cycle with an average 

period of about 11 years (Figure 15.21a). At the time of 

solar maximum,when sunspots are most numerous, we 

may see dozens of sunspots on the Sun at one time. In contrast, 

we see few if any sunspots at the time of solar minimum.The 

frequency of prominences and flares also follows 

the sunspot cycle, with these events being most common 

at solar maximum and least common at solar minimum. 

Although we’ll call it an “11-year” cycle, the interval 

between solar maxima is sometimes as long as 15 years 

or as short as 7 years. The number of sunspots also varies 

dramatically (Figure 15.21a). In fact, sunspot activity virtually 

ceased between the years 1645 and 1715, a period 

sometimes called the Maunder minimum (after E. W. 

Maunder, who identified it in historical sunspot records). 

Another feature of the sunspot cycle is a gradual 

change in the solar latitudes at which individual sunspots 

form and dissolve (Figure 15.21b). As a cycle begins at solar 

minimum, sunspots form primarily at mid-latitudes (30° 

to 40°) on the Sun. The sunspots tend to form at lower latitudes 

as the cycle progresses, appearing very close to the 

solar equator as the next solar minimum approaches. 

A less obvious feature of the sunspot cycle is that 

something peculiar happens to the Sun’s magnetic field 

at each solar minimum. The field lines connecting all pairs 

of sunspots (see Figure 15.17) tend to point in the same 

direction throughout an 11-year solar cycle (within each 

hemisphere). For example, all compass needles might point 

from the easternmost sunspot to the westernmost sunspot 

in a pair. However, as the cycle ends at solar minimum, the 

magnetic field reverses: In the subsequent solar cycle, the 

field lines connecting pairs of sunspots point in the opposite 

direction. Apparently, the entire magnetic field of the Sun 

flip-flops every 11 years. 

The magnetic reversals hint that the sunspot cycle 

is related to the generation of magnetic fields on the Sun. 

They also tell us that the complete magnetic cycle of the 

Sun, called the solar cycle, really averages 22 years, since it 

takes two 11-year cycles before the magnetic field is back 

the way it started. 


0.5% 

percentage of Sun’s surface 

covered by sunspots 

0.4% 

0.3% 

0.2% 

0.1% 

0.0% 

1900 1910 1920 1930 1940 1950 1960 1970 1980 1990 2000 

a This graph shows how the number of sunspots on the Sun changes with time. The vertical axis shows the 

percentage of the Sun’s surface covered by sunspots. The cycle has a period of approximately 11 years. 

year 

90 N 

30 N 

latitude 

Equator 

30 S 

90 S 

1900 

1910 1920 1930 1940 1950 1960 1970 1980 1990 2000 

b This graph shows how the latitudes at which sunspot groups appear tend to shift during a single sunspot 

cycle. Each dot represents a group of sunspots and indicates the year (horizontal axis) and latitude (vertical 

axis) at which the group appeared. 

Figure 15.21 Sunspot cycle during the past century. 

year 

What Causes the Sunspot Cycle? 

The causes of the Sun’s magnetic fields and the sunspot 

cycle are not well understood, but we believe we know 

the general nature of the processes involved. Convection 

is thought to dredge up weak magnetic fields generated in 

the solar interior, amplifying them as they rise. The Sun’s 

rotation—faster at its equator than near its poles—then 

stretches and shapes these fields. 

Imagine what happens to a magnetic field line that originally 

runs along the Sun’s surface directly from the north 

pole to the south pole. At the equator, the field line circles 

the Sun in 25 days, but at higher latitudes the field line lags 

behind. Gradually, this rotation pattern winds the field line 

more and more tightly around the Sun (Figure 15.22). This 

process, operating at all times over the entire Sun, produces 

the contorted field lines that generate sunspots and other 

solar activity. 

Investigating how the Sun’s magnetic field develops 

and changes in time requires sophisticated computer models. 

Scientists are working hard on such models, but the 

behavior of these fields is so complex that approximations 

are necessary even with the best supercomputers. Using 

these computer models, scientists have successfully replicated 

some features of the sunspot cycle, such as changes 

in the number and latitude of sunspots and the magnetic 

field reversals that occur about every 11 years. However, 

much still remains mysterious, including why the period 

of the sunspot cycle varies and why solar activity is different 

from one cycle to the next. 

Solar Activity and Earth 

During solar maximum, solar flares and other forms of 

solar activity send large numbers of highly energetic charged 

particles (protons and electrons) toward Earth. Sometimes 

these particles travel in the smooth flow known as the solar 

wind. Other times they come in the form of huge magnetic 

bubbles called coronal mass ejections. Do these forms of solar 

weather affect Earth? In at least some ways, the answer is 

a definitive yes. 


e quator 

N N N 

magnetic 

field line 

differing 

rotation 

rates 

time 

time 

S S S 

Figure 15.22 The Sun rotates more quickly at its equator than it does near its poles. Because gas circles 

the Sun faster at the equator, it drags the Sun’s north-south magnetic field lines into a more twisted configuration. 

The magnetic field lines linking pairs of sunspots, depicted here as green and black blobs, trace out 

the directions of these stretched and distorted field lines. 

The magnetic field associated with the solar wind constantly 

interacts with Earth’s magnetic field. Occasionally 

these fields interconnect. When that happens, large amounts 

of energy are released from the magnetic field into the 

charged particles near the interconnection zone. Many 

of these energized particles then flow down Earth’s magnetic 

field lines toward the poles (Figure 15.23a). Collisions 

between the charged particles and atoms in Earth’s upper 

atmosphere cause electrons in the atoms to jump to higher 

energy levels [Section 4.4].These excited atoms subsequently 

emit visible-light photons as they drop to lower energy levels, 

creating the shimmering light of auroras (Figure 15.23b). 

Because coronal mass ejections are particularly energetic, the 

auroras they stimulate can be especially spectacular. 

Particles streaming from the Sun after the occurrence 

of solar flares, coronal mass ejections, or other major solar 

storms can also have practical impacts on society. For example, 

these particles can hamper radio communications, 

disrupt electrical power delivery, and damage the electronic 

components in orbiting satellites. During a particularly 

powerful magnetic storm on the Sun in March 1989, the 

U.S. Air Force temporarily lost track of over 2,000 satellites, 

SPECIAL TOPIC 

Long-Term Change in Solar Activity 

Figure 15.21 shows that the sunspot cycle varies in length and 

intensity, and it sometimes seems to disappear altogether. With 

these facts as background, many scientists are searching for longerterm 

patterns in solar activity. Unfortunately, the search for longerterm 

variations is difficult because telescopic observations of 

sunspots cover a period of only about 400 years. Some naked-eye 

observations of sunspots recorded by Chinese astronomers go 

back almost 2,000 years, but these records are sparse, and nakedeye 

observations may not be very reliable. We can also guess at 

past solar activity from descriptions of solar eclipses recorded 

around the world: When the Sun is more active, the corona tends 

to have longer and brighter “streamers” visible to the naked eye. 

Another way to gauge past solar activity is to study the amount 

of carbon-14 in tree rings. High-energy cosmic rays [Section 19.2] 

coming from beyond our own solar system produce radioactive 

carbon-14 in Earth’s atmosphere. During periods of high solar 

activity, the solar wind tends to grow stronger, shielding Earth 

from some of these cosmic rays. Thus, production of carbon-14 

drops when the Sun is more active. All the while, trees steadily 

breathe in atmospheric carbon, in the form of carbon dioxide, 

and incorporate it year by year into each ring. We can therefore 

estimate the level of solar activity in any given year by measuring 

the level of carbon-14 in the corresponding ring. No clear evidence 

has yet been found of longer-term cycles of solar activity, 

but the search goes on. 

Theoretical models predict a very long term trend of lessening 

solar activity. According to our theory of solar system formation, 

the Sun must have rotated much faster when it was young 

[Section 9.3].Because a combination of convection and rotation 

generates solar activity, a faster rotation rate should have meant 

much more activity. Observations of other stars that are similar 

to the Sun but rotate faster confirm that these stars are much 

more active. We find evidence for many more “starspots” on these 

stars than on the Sun, and their relatively bright ultraviolet and 

X-ray emissions suggest that they have brighter chromospheres 

and coronas—just as we would expect if they are more active 

than the Sun. 

chapter 15 • Our Star 515

particles spiral around 

magnetic field lines 

stream of solar particles 

from solar wind 

N 

S 

a Interactions between Earth’s magnetic field and the magnetic field of the solar wind can 

send energetic charged particles toward Earth’s poles. 

Figure 15.23 Particles from the Sun cause auroras on Earth. 

b Aurora along the coast of Norway. 

and powerful currents induced in the ground circuits of 

the Quebec hydroelectric system caused it to collapse for 

more than 8 hours. The combined cost of the loss of power 

in the United States and Canada exceeded $100 million. In 

January 1997, AT&T lost contact with a $200-million communications 

satellite, probably because of damage caused 

by particles coming from another powerful solar storm. 

Satellites in low-Earth orbit are particularly vulnerable 

during solar maximum, when the increase in solar X rays 

and energetic particles heats Earth’s upper atmosphere, 

causing it to expand. The density of the gas surrounding 

low-flying satellites therefore rises, exerting drag that saps 

their energy and angular momentum. If this drag proceeds 

unchecked, the satellites ultimately plummet back to Earth. 

Satellites in low orbits, including the Hubble Space Telescope 

and the Space Station, require occasional boosts to prevent 

them from falling out of the sky. 

Connections between solar activity and Earth’s climate 

are much less clear. The period from 1645 to 1715, when 

solar activity seems to have virtually ceased, was a time 

of exceptionally low temperatures in Europe and North 

America known as the Little Ice Age. Did the low solar activity 

cause these low temperatures, or was their occurrence 

a coincidence? No one knows for sure. Some researchers 

have claimed that certain weather phenomena, such as 

drought cycles or frequencies of storms, are correlated with 

the 11- or 22-year cycle of solar activity. However, the data 

supporting these correlations are weak in many cases, and 

even real correlations may be coincidental. 

Part of the difficulty in linking solar activity with climate 

is that no one understands how the linkage might 

work.Although emissions of ultraviolet light, X rays, and 

high-energy particles increase substantially from solar minimum 

to solar maximum, the total luminosity of the Sun 

barely changes at all. (The Sun becomes only about 0.1% 

brighter during solar maximum.) Thus, if solar activity 

really is affecting Earth’s climate, it must be through some 

very subtle mechanism. For example, perhaps the expansion 

of Earth’s upper atmosphere that occurs with solar 

maximum somehow causes changes in weather. 

The question of how solar activity is linked to Earth’s 

climate is very important, because we need to know whether 

global warming is affected by solar activity in addition to 

human activity. Unfortunately, for the time being at least, 

we can say little about this question. 


THE BIG PICTURE 

Putting Chapter 15 into Context 

In this chapter, we have examined our Sun, the nearest 

star. When you look back at this chapter, make sure you 

understand these “big picture” ideas: 

● The ancient riddle of why the Sun shines is now solved. 

The Sun shines with energy generated by fusion of 

hydrogen into helium in the Sun’s core. After a millionyear 

journey through the solar interior and an 8-minute 

journey through space, a small fraction of this energy 

reaches Earth and supplies sunlight and heat. 

● Gravitational equilibrium, the balance between pressure 

and gravity, determines the Sun’s interior structure 

and maintains its steady nuclear burning rate. 

The Sun achieved its long-lasting state of gravitational 

equilibrium when energy generation by fusion in the 

core came into balance with the energy lost through 

thermal radiation from the surface. If the Sun were not 

relatively steady, life on Earth might not have been 

possible. 

● The Sun’s atmosphere displays its own version of 

weather and climate, governed by solar magnetic fields. 

Solar weather has important influences on Earth. 

● The Sun is important not only as our source of light 

and heat, but also because it is the only star near 

enough for us to study in great detail. In the coming 

chapters, we will use what we’ve learned about the 

Sun to help us understand other stars. 

SUMMARY OF KEY CONCEPTS 


• What process creates energy in the Sun? Fusion of 

hydrogen into helium in the Sun’s core generates the 

Sun’s energy. 

• Why does the Sun’s size remain stable? The Sun’s size 

remains stable because it is in gravitational equilibrium. 

The outward pressure of hot gas balances the 

inward force of gravity at every point within the Sun. 

• How did the Sun become hot enough for fusion in the 

first place? As the Sun was forming, it grew hotter 

as it shrank in size because gravitational contraction 

converted gravitational potential energy into thermal 

energy. Gravitational contraction continued to 

shrink the Sun and raise its central temperature until 

the core became hot and dense enough for nuclear 

fusion. 

15.2 Plunging to the Center of the Sun: 


• What are the major layers of the Sun, from the center 

out? The layers of the Sun are core, radiation zone, 

convection zone, photosphere, chromosphere, and 

corona. 

• What do we mean by the “surface” of the Sun? We 

consider the photosphere to be the surface of the Sun 

because light can pass through the photosphere but 

cannot escape from deeper inside the Sun. Thus, 

photographs of visible light from the Sun show us 

what the photosphere looks like. 

• What is the Sun made of ? It is made almost entirely 

of hydrogen and helium (98% of the Sun’s mass). 


• Why does fusion occur in the Sun’s core? The core 

temperature and pressure are so high that colliding 

nuclei can come close enough together for the strong 

force to overcome electromagnetic repulsion and 

bind them together. 

• Why is energy produced in the Sun at such a steady 

rate? The fusion rate is self-regulating like a thermostat. 

If the fusion rate increases for some reason, the 

added energy production puffs up and cools the core, 

bringing the rate back down. Similarly, a decrease 

in the fusion rate allows the core to shrink and heat, 

bringing the fusion rate back up. 

• Why was the Sun dimmer in the distant past? 

Although the fusion rate is steady on short time 

scales, it gradually increases over billions of years, 

increasing the Sun’s luminosity. The increase occurs 

because fusion gradually reduces the number of individual 

nuclei in the solar core. Four hydrogen nuclei 

are fused to make just one helium nucleus, causing 

the core to shrink and become hotter. 

• How do we know what is happening inside the Sun? 

We can construct theoretical models of the solar 

interior using known laws of physics and then check 

continued 


the models against observations of the Sun’s size, 

surface temperature, and energy output as well as 

studies of “sun quakes” and solar neutrinos. 

• What is the solar neutrino problem? Is it solved? 

Neutrino detectors capture fewer neutrinos coming 

from the Sun than models of fusion in the core predict. 

This discrepancy is called the solar neutrino 

problem. The problem now appears to be solved. 

Apparently, neutrinos can transform themselves 

among three different types as they travel from the 

solar core to Earth, while most detectors can capture 

only one type. Thus, the detectors capture fewer 

than the expected number of neutrinos. 


• How long ago did fusion generate the energy we now 

receive as sunlight? Fusion created the energy we receive 

today about a million years ago. It takes about 

a million years for photons and then convection to 

transport energy through the solar interior to the 

photosphere. Once sunlight emerges from the photosphere, 

it takes only about 8 minutes to reach Earth. 

• How are sunspots, prominences, and flares related to 

magnetic fields? Sunspots occur where strong magnetic 

fields trap and isolate gas from the surrounding 

plasma of the photosphere. The trapped gas cools, 

so the sunspots become cooler and darker than the 

rest of the photosphere. Sunspots tend to occur in 

pairs connected by a loop of magnetic field, which 

may rise high above the surface as a solar prominence. 

The magnetic fields are twisted and contorted 

by the Sun’s rotation, and solar flares may occur 

when the field lines suddenly snap and release their 

energy. 

• What is surprising about the temperature of the 

chromosphere and corona, and how do we explain it? 

Temperature gradually decreases from the core to 

the photosphere but then rises again in the chromosphere 

and corona. These high layers of the Sun are 

probably heated by energy carried upward along the 

magnetic field lines by waves that are generated as 

turbulent motions in the convection zone shake the 

magnetic field lines. 

15.5 Solar Weather and Climate 

• What is the sunspot cycle? The sunspot cycle, or the 

variation in the number of sunspots on the Sun’s 

surface, has an average period of 11 years. The magnetic 

field flip-flops every 11 years or so, resulting 

in a 22-year magnetic cycle. Sunspots first appear 

at mid-latitudes at solar minimum, then become 

increasingly more common near the Sun’s equator 

as the next minimum approaches. 

• What effect does solar activity have on Earth and its 

inhabitants? Particles ejected from the Sun by solar 

flares and other types of solar activity can affect 

communications, electrical power delivery, and the 

electronic circuits in space vehicles. The connections 

between solar activity and Earth’s climate are not 

clear. 

Sensible Statements? 

Decide whether each of the following statements is sensible and 

explain why it is or is not. 

1. Before Einstein, gravitational contraction appeared to be a 

perfectly plausible mechanism for solar energy generation. 

2. A sudden temperature rise in the Sun’s core is nothing to 

worry about, because conditions in the core will soon return 

to normal. 

3. If fusion in the solar core ceased today, worldwide panic 

would break out tomorrow as the Sun began to grow 

dimmer. 

4. Astronomers have recently photographed magnetic fields 

churning deep beneath the solar photosphere. 

5. Neutrinos probably can’t harm me, but just to be safe 

I think I’ll wear a lead vest. 

6. If you want to see lots of sunspots, just wait for solar 

maximum! 

7. News of a major solar flare today caused concern among 

professionals in the fields of communications and electrical 

power generation. 

8. By observing solar neutrinos, we can learn about nuclear 

fusion deep in the Sun’s core. 


Problems 

9. Gravitational Contraction. Briefly describe how gravitational 

contraction generates energy and when it was important 

in the Sun’s history. 

10. Solar Characteristics. Briefly describe the Sun’s luminosity, 

mass, radius, and average surface temperature. 

11. Sunspots. What are sunspots? Why do they appear dark in 

pictures of the Sun? 

12. Solar Fusion. What is the overall nuclear fusion reaction in 

the Sun? Briefly describe the proton–proton chain. 

13. Models of the Sun. Explain how mathematical models allow 

us to predict conditions inside the Sun. How can we be 

confident that the models are on the right track? 

14. Sun Quakes. How are “sun quakes” similar to earthquakes? 

How are they different? Describe how we can observe them 

and how they help us learn about the solar interior. 

15. Energy Transport. Why does the energy produced by fusion 

in the solar core take so long to reach the solar surface? 

Describe the processes of radiative diffusion and convection 

in the solar interior. 

16. The Photosphere. Describe the appearance and temperature 

of the Sun’s photosphere. What is granulation? How would 

granulation appear in a movie? 

17. Observing the Sun’s Atmosphere. Why is the chromosphere 

best viewed with ultraviolet telescopes? Why is the corona 

best viewed with X-ray telescopes? 

18. An Angry Sun. A Time magazine cover once suggested that 

an “angry Sun” was becoming more active as human activity 

changed Earth’s climate through global warming. It’s 

certainly possible for the Sun to become more active at the 

same time that humans are affecting Earth, but is it possible 

that the Sun could be responding to human activity? Can 

humans affect the Sun in any significant way? Explain. 

*19. Number of Fusion Reactions in the Sun. Use the fact that each 

cycle of the proton–proton chain converts 4.7 10 29 kg 

of mass into energy (see Mathematical Insight 15.1), along 

with the fact that the Sun loses a total of about 4.2 10 9 kg 

of mass each second, to calculate the total number of times 

the proton–proton chain occurs each second in the Sun. 

*20. The Lifetime of the Sun. The total mass of the Sun is about 

2 10 30 kg, of which about 75% was hydrogen when the 

Sun formed. However, only about 13% of this hydrogen 

ever becomes available for fusion in the core. The rest remains 

in layers of the Sun where the temperature is too low 

for fusion. 

a. Based on the given information, calculate the total 

mass of hydrogen available for fusion over the lifetime 


b. Combine your results from part (a) and the fact that the 

Sun fuses about 600 billion kg of hydrogen each second 

to calculate how long the Sun’s initial supply of hydrogen 

can last. Give your answer in both seconds and years. 

c. Given that our solar system is now about 4.6 billion years 

old, when will we need to worry about the Sun running 

out of hydrogen for fusion? 

*21. Solar Power Collectors. This problem leads you through the 

calculation and discussion of how much solar power can be 

collected by solar cells on Earth. 

a. Imagine a giant sphere surrounding the Sun with a 

radius of 1 AU. What is the surface area of this sphere, 

in square meters? (Hint: The formula for the surface 

area of a sphere is 4pr 2 .) 

b. Because this imaginary giant sphere surrounds the Sun, 

the Sun’s entire luminosity of 3.8 10 26 watts must 

pass through it. Calculate the power passing through 

each square meter of this imaginary sphere in watts per 

square meter. Explain why this number represents the 

maximum power per square meter that can be collected 

by a solar collector in Earth orbit. 

c. List several reasons why the average power per square 

meter collected by a solar collector on the ground will 

always be less than what you found in part (b). 

d. Suppose you want to put a solar collector on your roof. 

If you want to optimize the amount of power you can 

collect, how should you orient the collector? (Hint: The 

optimum orientation depends on both your latitude 

and the time of year and day.) 

*22. Solar Power for the United States. The total annual U.S. energy 

consumption is about 2 10 20 joules. 

a. What is the average power requirement for the United 

States, in watts? (Hint: 1 watt 1 joule/s.) 

b. With current technologies and solar collectors on the 

ground, the best we can hope is that solar cells will 

generate an average (day and night) power of about 

200 watts/m 2 .(You might compare this to the maximum 

power per square meter you found in problem 22b.) 

What total area would we need to cover with solar cells 

to supply all the power needed for the United States? 

Give your answer in both square meters and square 

kilometers. 

c. The total surface area of the United States is about 

2 10 7 km 2 .What fraction of the U.S. area would have 

to be covered by solar collectors to generate all of the 

U.S. power needs? In one page or less, describe potential 

environmental impacts of covering so much area with 

solar collectors. Also discuss whether you think these 

environmental impacts would be greater or less than 

the impacts of using current energy sources such as coal, 

oil, nuclear power, and hydroelectric power. 


*23. The Color of the Sun. The Sun’s average surface temperature 

is about 5,800 K. Use Wien’s law (see Mathematical Insight 

6.2) to calculate the wavelength of peak thermal emission 

from the Sun. What color does this wavelength correspond 

to in the visible-light spectrum? In light of your answer, why 

do you think the Sun appears white or yellow to our eyes? 

Discussion Questions 

24. The Role of the Sun. Briefly discuss how the Sun affects us 

here on Earth. Be sure to consider not only factors such as 

its light and warmth, but also how the study of the Sun has 

led us to new understanding in science and to technological 

developments. Overall, how important has solar research 

been to our lives? 

25. The Solar Neutrino Problem. Discuss the solar neutrino problem 

and its potential solutions. How serious do you consider 

this problem? Do you think current theoretical models 

of the Sun could be wrong in any fundamental way? 

Why or why not? 

26. The Sun and Global Warming. One of the most pressing 

environmental issues on Earth concerns the extent to which 

human emissions of greenhouse gases are warming our 

planet. Some people claim that part or all of the observed 

warming over the past century may be due to changes on 

the Sun, rather than to anything humans have done. Discuss 

how a better understanding of the Sun might help us 

understand the threat posed by greenhouse gas emissions. 

Why is it so difficult to develop a clear understanding of 

how the Sun affects Earth’s climate? 

Web Projects 

MEDIA EXPLORATIONS 

Take advantage of the useful Web links on www. 

astronomyplace.com to assist you with the following 

projects. 

1. Current Solar Activity. Daily information about solar 

activity is available at NASA’s Web site sunspotcycle. 

com. Where are we in the sunspot cycle right now? 

When is the next solar maximum or minimum expected? 

Have there been any major solar storms in 

the past few months? If so, did they have any significant 

effects on Earth? Summarize your findings in 

a one- to two-page report. 

2. Solar Observatories in Space. Visit NASA’s Web site 

for the Sun–Earth connection and explore some 

of the current and planned space missions designed 

to observe the Sun. Choose one mission to study in 

greater depth, and write a one- to two-page report 

on the mission status and goals and what it has 

taught or will teach us about the Sun. 

3. Sudbury Neutrino Observatory. Visit the Web site for 

the Sudbury Neutrino Observatory (SNO) and learn 

how it has helped to solve the solar neutrino problem. 

Write a one- to two-page report describing the 

observatory, any recent results, and what we can 

expect from it in the future.

Chapter 15--Our Sun - Geological Sciences

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