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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?
Answers appear at the end of the chapter.
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?
Answers appear at the end of the chapter.
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)
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
Distance from Sun’s center (solar radii)
Distance from Sun’s center (solar radii)
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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Probing the Dynamic Sun 21
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?
Answers appear at the end of the chapter.
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?
Answer appears at the end of the chapter.
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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Probing the Dynamic Sun 21
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?
Answer appears at the end of the chapter.
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?
Answer appears at the end of the chapter.
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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Probing the Dynamic Sun 21
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?
Answer appears at the end of the chapter.
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?
Answer appears at the end of the chapter.
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?
Answer appears at the end of the chapter.
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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Probing the Dynamic Sun 22
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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Probing the Dynamic Sun 22
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?
Answer appears at the end of the chapter.
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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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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Probing the Dynamic Sun 22
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?
Answer appears at the end of the chapter.
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?
Answer appears at the end of the chapter.
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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Probing the Dynamic Sun 22
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?
Answer appears at the end of the chapter.
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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Probing the Dynamic Sun 22
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