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Chapter 12
How the Stars Shine:
Cosmic Furnaces
Introduction
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Even though individual stars shine for a
relatively long time, they are not eternal.
Stars are born out of the gas and dust
that exist within a galaxy; they then begin
to shine brightly on their own.
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Eventually, they die.
Though we can directly observe only the
outer layers of stars, we can deduce that
the temperatures at their centers must be
millions of kelvins.
We can even figure out what it is deep
down inside that makes the stars shine.
To determine the probable life history of a
typical star, we observe stars having
many different ages and assume that
they evolve in a similar manner.
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However, we must take into account the
different masses of stars; some aspects of
their evolution depend critically on mass.
Introduction
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We start this chapter by discussing the birth of
stars.
We see how new capabilities of observing in the
infrared in addition to the visible are helping us
understand star formation (see figure).
We then consider the processes that go on inside
a star during its life on the main sequence.
Finally, we begin the story of the evolution of
stars when they finish the main-sequence stage
of their lives.
Chapters 13 and 14 will continue the story of
what is called “stellar evolution,” all the way to
the deaths of stars.
Introduction
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Near the end of this chapter we will see that the most
important experiment to test whether we understand how
stars shine is the search for elusive particles, called
neutrinos, from the Sun.
Over the past decades a search for them has been made,
but only about a third to half of those expected had been
found.
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Recent experiments have provided better ways of detecting
neutrinos than we previously had, and they were there all
along, though transformed and thus hidden!
The results indicate that we did not understand neutrinos
as well as we had thought.
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These astronomical results therefore have added important
knowledge about fundamental physics in addition to our
understanding of the stars.
12.1 Starbirth
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The birth of a star begins with a nebula—a
large region of gas and dust (see figures).
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The dust (tiny solid particles) may have
escaped from the outer atmospheres of
giant stars.
12.1 Starbirth
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The regions of gas and dust (often called clouds, or “giant
molecular clouds”) from which stars are forming are best
observed in the infrared and radio regions of the spectrum,
because most other forms of radiation (such as optical and
ultraviolet) cannot penetrate them.
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We discuss the infrared observations largely here, including the
new capabilities of NASA’s Spitzer Space Telescope, and we leave
the radio observations to Chapter 15 on the Milky Way Galaxy.
Stars forming at the present time incorporate this previously
cycled gas and dust, which gives them their relatively high
abundances of elements heavier than helium in the periodic
table (still totaling less than a per cent).
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In contrast, the oldest stars we see formed long ago when only
primordial hydrogen and helium were present, and therefore have
lower abundances of the heavier elements, as we discussed near
the end of the previous chapter.
12.1a Collapse of a Cloud
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Consider a region that reaches a higher density than its
surroundings, perhaps from a random fluctuation in
density or—in a leading theory of why galaxies have spiral
arms—because a wave of compression passes by.
Still another possibility is that a nearby star explodes (a
“supernova”; see Chapter 13), sending out a shock wave
that compresses the gas and dust.
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In any case, once the cloud gains a higher-than-average
density, the gas and dust continue to collapse due to gravity.
Energy is released, and the material accelerates inward.
Magnetic fields may resist the infalling gas, slowing the
infall, though the role of magnetic fields is not well
understood in detail.
12.1a Collapse of a Cloud
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Eventually, dense cores, each with a mass
comparable to that of a star, form and grow like tiny
seeds within the vast cloud.
These protostars (from the prefix of Greek origin
meaning “primitive”), which will collectively form a
star cluster, continue to collapse, almost unopposed
by internal pressure.
But when a protostar becomes sufficiently dense,
frequent collisions occur among its particles; hence,
part of the gravitational energy released during
subsequent collapse goes into heating the gas,
increasing its internal pressure. (In general,
compression heats a gas; for example, a bicycle tire
feels warm after air is vigorously pumped into it.)
12.1a Collapse of a Cloud
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The rising internal pressure, which is highest in the
protostar’s center and decreases outward, slows down the
collapse until it becomes very gradual and more accurately
described as contraction.
The object is now called a pre-main-sequence star
(see figure).
12.1a Collapse of a Cloud
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By this time, the object has contracted by a huge
fraction, a factor of 10 million, from about 10
trillion km across to about a million km across—
that is, something initially larger than the whole
Solar System collapses until most of its mass is in
the form of a single star.
During the contraction phase, a disk tends to form
because the original nebula was rotating slightly.
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We discussed this process when considering the
“nebular hypothesis” for the formation of our Solar
System (Chapter 9).
Dusty disks have been found around young stars
and pre-main-sequence stars in nebulae such as
the Orion Nebula (see figure).
These are sometimes called protoplanetary disks
(“proplyds”), and they support the theoretical
expectation that planetary systems are common.
12.1a Collapse of a Cloud
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Jets of gas are commonly ejected in opposite directions
out the poles of the rotating pre-mainsequence star (see
figure, right).
As energy is radiated from the surface of the pre-mainsequence star, its internal pressure decreases, and it
gradually contracts.
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This release of gravitational energy heats the interior,
thereby increasing the internal temperature and pressure.
It is also the source of the radiated energy.
Gravitational energy was released in this way in the early
Solar System.
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As the temperature in the interior rises, the
outward force resulting from the outwardly
decreasing pressure increases, and
eventually it balances the inward force of
gravity, a condition known as “hydrostatic
equilibrium.”
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As we shall discuss later, this mechanical
balance is the key to understanding stable
stars; see Figure at left.
12.1a Collapse of a Cloud
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Theoretical analysis shows that the dust surrounding the
stellar embryo we call a pre-main-sequence star should
absorb much of the radiation that the object emits.
The radiation from the pre-main-sequence star should
heat the dust to temperatures that cause it to radiate
primarily in the infrared.
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Infrared astronomers have found many objects that are
especially bright in the infrared but that have no known
optical counterparts.
These objects seem to be located in regions where the
presence of a lot of dust, gas, and young stars indicates
that star formation might still be going on.
12.1a Collapse of a Cloud
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Imaging in the visible (with the Hubble Space Telescope) and in the
infrared (not only with previous infrared space telescopes but now
especially with the Spitzer Space Telescope) has shown how young stars
are born inside giant pillars of gas and dust inside certain nebulae.
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The Eagle Nebula is the most famous example because of the beautiful
Hubble image showing exquisite detail, with false colors assigned to different
filters (see figure).
12.1a Collapse of a Cloud
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As hot stars form, their intense radiation
evaporates the gas and dust around them,
freeing them from the cocoons of gas and
dust in which they were born.
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We see this “evaporation” taking place at the
tops of the Eagle Nebula’s “pillars.”
The stars are destroying their birthplaces as
they become independent and more visible
from afar.
12.1b The Birth Cries of Stars
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To their surprise, astronomers have discovered that young stars
send matter out in oppositely directed beams, while they had
expected to find only evidence of infall.
This “bipolar ejection” (see figure) may imply that a disk of matter
orbits such premain-sequence stars, blocking an outward flow of
gas in the equatorial direction and later coalescing into planets.
Thus the flow of gas is channeled toward the poles.
12.1b The Birth Cries of Stars
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Sometimes clumps of gas appear, but only recently have
they been identified with ejections from stars in the
process of collapsing.
The clumps seem like spinning bullets, though what
makes them spin is uncertain (perhaps connected with the
magnetic field).
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The ejection of these spinning clumps helps slow the star’s
rate of spin, since they carry away angular momentum from
what had been a rapidly spinning pre-mainsequence star.
At the same time, some gas with low angular momentum
is falling in toward the star.
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Hidden here is perhaps the main unsolved problem in star
formation at the moment: How do stars figure out what
their final mass will be?
12.1b The Birth Cries of Stars
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The bipolar ejection appears as “Herbig-Haro objects,”
clouds of interstellar gas heated by shock waves from jets
of high-speed gas.
The jets are being ejected from a premain-sequence star,
a star in the process of being born.
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Since the pre-mainsequence star is
hidden in visible light
by a dusty cocoon of
gas, infrared
observations of
HerbigHaro objects most
clearly reveal what is
going on (see figure).
12.1b The Birth Cries of Stars
The jets of gas were formed as the pre-main-sequence star contracted
under the force of its own gravity.
Because a thick disk of cool gas and dust surrounds the premainsequence star, the gas squirts outward along the pre-main-sequence
star’s axis of rotation at speeds of perhaps 1 million km /hr.
HH–1 and HH–2 (see figure) are more irregular in shape than many
other Herbig-Haro objects, perhaps because the bow shock wave we are
seeing (a shock wave like those formed by the bow of a boat plowing
through the water) has broken up.
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These objects are about 1500
light-years from us, in a starforming region of the
constellation Orion.
The smallest features resolved
are about the size of our Solar
System, and the whole image is
only about 1 light-year across.
12.1b The Birth Cries of Stars
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Several classes of stars that vary erratically in brightness have
been found.
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One of these classes, called T Tauri, contains pre-main-sequence
stars as massive as or less massive than the Sun.
Presumably, these stars are so young that they have not quite
settled down to a steady and reliable existence on the main
sequence. (T Tauri stars always have the word “stars” in their
name though technically they haven’t reached the main
sequence, so they are not yet fully formed stars.)
In astronomical teaching, we have the question of whether to
first consider the formation of stars in the star section of the
book, as here, or in the section about the gas and dust between
the stars from which the stars form.
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We choose to do some of each, and will continue our discussion of
stars in formation in that latter location, Chapter 15 on the Milky
Way Galaxy.
12.2 Where Stars Get Their Energy
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If the Sun got all of its energy from gravitational
contraction, it could have shined for only about 30 million
years, not very long on an astronomical timescale.
Yet we know that rocks about 4 billion years old have
been found on Earth, and up to 4.4 billion years old on
the Moon, so the Sun and the Solar System have been
around at least that long.
Moreover, fossil records of planets and animals, which
presumably used the Sun’s light and heat, date back
billions of years.
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Some other source of energy must hold the Sun and other
stars up against their own gravitational pull.
12.2 Where Stars Get Their Energy
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Actually, a pre-main-sequence star will heat up until its
central portions become hot enough (at least one million
kelvins) for nuclear fusion to take place, at which time it
reaches the main sequence of the temperature-luminosity
(temperature-magnitude, or Hertzsprung-Russell; see
Chapter 11) diagram.
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Using this process, which we will soon discuss in detail, the
star can generate enough energy to support it during its
entire lifetime on the main sequence.
A star’s luminosity and temperature change little while it is
on the main sequence; nuclear reactions provide the
stability.
12.2 Where Stars Get Their Energy
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The energy makes the
particles in the star move
around rapidly.
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Such rapid, random
motions in a gas are the
definition of high
temperature.
The thermal pressure, the force from these moving particles
pushing on each area of gas, is also high.
The varying pressure, which decreases outward from the center,
produces a force that pushes outward on any given pocket of gas.
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This outward force balances gravity’s inward pull on the pocket
(“hydrostatic equilibrium,” which we illustrated in the figure).
12.2 Where Stars Get Their Energy
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The basic fusion process in main-sequence stars fuses four hydrogen
nuclei into one helium nucleus.
In the process, tremendous amounts of energy are released. (Hydrogen
bombs on Earth fuse hydrogen nuclei into helium, but use different fusion
sequences. The fusion sequences that occur in stars are far too slow for
bombs.)
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A hydrogen nucleus is but a single proton.
A helium nucleus is more complex; it consists of two
protons and two neutrons (see figure).
The mass of the helium nucleus that is the final product
of the fusion process is slightly less than the sum of the
masses of the four hydrogen nuclei (protons) that went
into it.
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A small amount of the mass, m, “disappears” in the
process: 0.007 (0.7 per cent) of the mass of the four
protons.
12.2 Where Stars Get Their Energy
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The mass difference does not really disappear, but rather is
converted into energy, E, according to Albert Einstein’s famous
formula
E =mc2,
where c is the speed of light.
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This energy is known as the “binding energy” of the nucleus,
here specifically that of helium.
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Even though m is only a small fraction of the original mass, the
amount of energy released is prodigious; in the formula, c is a very
large number.
The loss of only 0.7 per cent of the central part of the Sun, for
example, is enough to allow the Sun to radiate as much as it does
at its present rate for a period of about ten billion (1010) years.
This fact, not realized until 1920 and worked out in more detail
in the 1930s, solved the longstanding problem of where the Sun
and the other stars get their energy.
12.2 Where Stars Get Their Energy
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All the main-sequence stars are approximately 90 per
cent hydrogen (that is, 90 per cent of the atoms are
hydrogen), so there is a lot of raw material to fuel the
nuclear “fires.”
We speak colloquially of “nuclear burning,” although,
of course, the processes are quite different from the
chemical processes that are involved in the “burning”
of logs or of autumn leaves.
In order to be able to discuss these processes, we
must first review the general structure of nuclei and
atoms.
12.3 Atoms and Nuclei
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As we mentioned in Chapter 2, an atom
consists of a small nucleus surrounded by
electrons.
Most of the mass of the atom is in the
nucleus, which takes up a very small
volume in the center of the atom.
The effective size of the atom, the
chemical interactions of atoms to form
molecules, and the nature of spectra are
all determined by the electrons.
12.3a Subatomic Particles
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The nuclear particles with which we need to be most familiar are the
proton and neutron.
Both of these particles have nearly the same mass, 1836 times greater
than the mass of an electron, though still tiny.
The neutron has no electric charge and the proton has one unit of
positive electric charge.
The electrons, which surround the nucleus, have one unit each of
negative electric charge.
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When an atom loses an electron, it has a net positive charge of 1 unit for
each electron lost.
The atom is now a form of ion (see figure).
12.3a Subatomic Particles
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Since the number of protons in the nucleus
determines the charge of the nucleus, it also dictates
the quota of electrons that the neutral state of the
atom must have.
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To be neutral, after all, there must be equal numbers of
positive and negative charges.
Each element (sometimes called “chemical
element”) is defined by the specific number of
protons in its nucleus.
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The element with one proton is hydrogen, that with
two protons is helium, that with three protons is
lithium, and so on.
12.3b Isotopes
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Though a given element always has the same number of protons in a
nucleus, it can have several different numbers of neutrons. (The number
of neutrons is usually somewhere between 1 and 2 times the number of
protons. The most common form of hydrogen, just a single proton, is the
main exception to this rule.)
The possible forms of the same element having different numbers of
neutrons are called isotopes.
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For example, the nucleus of ordinary hydrogen contains one proton and no
neutrons.
An isotope of hydrogen (see figure) called deuterium (and sometimes
“heavy hydrogen”) has one proton and one neutron.
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Another isotope of hydrogen called tritium has one proton and two neutrons.
12.3b Isotopes
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Most isotopes do not have specific names, and we keep track of the
numbers of protons and neutrons with a system of superscripts and
subscripts.
The subscript before the symbol denoting the element is the number of
protons (called the atomic number), and a superscript after the symbol is
the total number of protons and neutrons together (called the mass
number, or atomic mass).
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For example, 1H2 is deuterium, since deuterium has one proton, which gives
the subscript, and an atomic mass of 2, which gives the superscript. (Note
that 21H is also correct notation.)
Deuterium has atomic number equal to 1 and mass number equal to 2.
Similarly, 92U238 is an isotope of uranium with 92 protons (atomic number 92)
and mass number of 238, which is divided into 92 protons and 23892=146
neutrons.
Each element has only certain isotopes.
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For example, most of the naturally occurring helium is in the form 2He4, with
a much lesser amount as 2He3.
12.3c Radioactivity and Neutrinos
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Sometimes an isotope is not stable, in that after a time it
will spontaneously change into another isotope or
element; we say that such an isotope is radioactive.
The most massive elements, those past uranium, are all
radioactive, and have average lifetimes that are very
short.
It has been theoretically predicted that around element
114, elements should begin being somewhat more stable
again.
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The handful of atoms of element 114 and 116 discovered in
1998 and 1999 are more stable than those of slightly lower
mass numbers— lasting even about 5 seconds instead of a
small fraction of a second. (A claim that element 118 was
also discovered has been withdrawn.)
12.3c Radioactivity and Neutrinos
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During certain types of radioactive decay, as well as when
a free proton and electron combine to form a neutron, a
particle called a neutrino is given off.
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A neutrino is a neutral particle (its name comes from the
Italian for “little neutral one”).
Neutrinos have a very useful property for the purpose of
astronomy: They rarely interact at all with matter.
Thus when a neutrino is formed deep inside a star, it can
usually escape to the outside without interacting with any
of the matter in the star.
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A photon of electromagnetic radiation, on the other hand,
can travel only about 0.5 mm (on average) in a stellar
interior before it is absorbed, and it takes about a hundred
thousand years for a photon to zig and zag its way to the
surface.
12.3c Radioactivity and Neutrinos
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The elusiveness of the neutrino not only
makes it a valuable messenger—indeed, the
only possible direct messenger—carrying news
of the conditions inside the Sun at the present
time, but also makes it very difficult for us to
detect on Earth.
A careful experiment carried out over many
years has found only about ⅓ the expected
number of neutrinos, as we shall soon see.
12.4 Stars Shining Brightly
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Let us now use our knowledge of atomic nuclei to explain how stars
shine.
For a premain-sequence star, the energy from the gravitational
contraction goes into giving the individual particles greater speeds; that
is, the gas temperature rises.
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When atoms collide at high temperature, electrons get knocked away from
their nuclei, and the atoms become fully ionized.
The electrons and nuclei can move freely and separately in this “plasma.”
For nuclear fusion to begin, atomic nuclei must get close enough to
each other so that the force that holds nuclei together, the “strong
nuclear force” (to be discussed in Chapter 19), can play its part.
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But all nuclei have positive charges, because they are composed of protons
(which bear positive charges) and neutrons (which are neutral).
The positive charges on any two nuclei cause an electrical repulsion
between them, which tends to prevent fusion from taking place.
12.4 Stars Shining Brightly
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However, at the high temperatures (millions of kelvins)
typical of a stellar interior, some nuclei occasionally have
enough energy to overcome this electrical repulsion.
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They come sufficiently close to each other that they
essentially collide, and the strong nuclear force takes over.
Fusion on the main sequence proceeds in one of two
ways, as will be discussed below.
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Once nuclear fusion begins, enough energy is generated to
maintain the pressure and prevent further contraction.
The pressure provides a force that pushes outward strongly
enough to balance gravity’s inward pull.
12.4 Stars Shining Brightly
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In the center of a star, the fusion process is selfregulating.
The star finds a balance between thermal pressure
pushing out and gravity pushing in.
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It thus achieves stability on the main sequence (at a
constant temperature and luminosity).
When we learn how to control fusion in power-generating
stations on Earth, which currently seems decades off (and
has long seemed so), our energy crisis will be over, since
deuterium, the potential “fuel,” is so abundant in Earth’s
oceans.
12.4 Stars Shining Brightly
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The greater a star’s mass, the hotter its core
becomes before it generates enough pressure to
counteract gravity.
The hotter core gives off more energy, so the
star becomes brighter (see figure), explaining
why main-sequence stars of large mass have
high luminosity.
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In fact, it turns out that more massive stars use
their nuclear fuel at a very much higher rate than
less massive stars.
Even though the more massive stars have more fuel to burn, they go
through it relatively quickly and live shorter lives than low-mass stars, as
we discussed in Chapter 11.
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The next two chapters examine the ultimate fates of stars, with the fates
differing depending on the masses of the stars.
12.5 Why Stars Shine
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Several chains of nuclear reactions have been
proposed to account for the fusion of four
hydrogen nuclei into a single helium nucleus.
Hans Bethe of Cornell University suggested some
of these procedures during the 1930s.
The different chain reactions prevail at different
temperatures, so chains that are dominant in
very hot stars may be different from the ones in
cooler stars.
12.5 Why Stars Shine
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When the temperature of the center of a main-sequence star is less than
about 20 million kelvins, the proton-proton chain (see figure)
dominates.
This sequence uses six hydrogen nuclei (protons), and winds up with one
helium nucleus plus two protons.
The net transformation is four hydrogen nuclei into one helium nucleus.
(Though two of the protons turn into neutrons, here this isn’t the main
point.)
12.5 Why Stars Shine
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But the original six protons contained more mass than do the
final single helium nucleus plus two protons.
The small fraction of mass that disappears is converted into an
amount of energy that we can calculate with the formula
E=mc2.
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For stellar interiors significantly hotter than that of the Sun, the
carbon-nitrogen oxygen (CNO) cycle dominates.
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According to Einstein’s special theory of relativity, mass and energy
are equivalent and interchangeable, linked by this equation.
This cycle begins with the fusion of a hydrogen nucleus (proton)
with a carbon nucleus.
After many steps, and the insertion of four protons, we are left
with one helium nucleus plus a carbon nucleus.
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Thus, as much carbon remains at the end as there was at the
beginning, and the carbon can start the cycle again.
12.5 Why Stars Shine
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As in the proton-proton chain, four hydrogen nuclei have
been converted into one helium nucleus during the CNO
cycle, 0.7 per cent of the mass has been transformed, and
an equivalent amount of energy has been released
according to E=mc2.
Main-sequence stars more massive than about 1.1 times
the Sun are dominated by the CNO cycle.
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Later in their lives, when they are no longer on the main
sequence, stars can have even higher interior temperatures,
above 108 K.
They then fuse helium nuclei to make carbon nuclei.
The nucleus of a helium atom is called an “alpha particle”
for historical reasons.
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Since three helium nuclei (2He4) go into making a single
carbon nucleus (6C12), the procedure is known as the triplealpha process.
12.5 Why Stars Shine
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A series of other processes can build still heavier elements
inside very massive stars.
These processes, and other element-building methods,
are called nucleosynthesis (new´clee-oh-sin´tha-sis).
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The theory of nucleosynthesis in stars can account for the
abundances (proportions) we observe of the elements
heavier than helium.
Currently, we think that the synthesis of isotopes of the
lightest elements (hydrogen, helium, and lithium) took
place in the first few minutes after the origin of the
Universe (Chapter 19), though some of the observed
helium was produced later by stars.
12.6 Brown Dwarfs
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When a pre-main-sequence star has at least 7.5 per cent
of the Sun’s mass (that is, it has about 75 Jupiter
masses), nuclear reactions begin and continue, and it
becomes a normal star.
But if the mass is less than 7.5 per cent of the Sun’s mass,
the central temperature does not become hot enough for
nuclear reactions using ordinary hydrogen (protons) to be
sustained. (Masses of this size do, however, fuse
deuterium into helium, but this phase of nuclear fusion
doesn’t last long because there is so little deuterium in the
Universe relative to ordinary hydrogen.)
12.6 Brown Dwarfs
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These objects shine dimly, shrinking and dimming as
they age.
They came to be called brown dwarfs, mainly
because “brown” is a mixture of many colors and
people didn’t agree how such supposedly “failed stars”
would look, and also because they emit very little light
(see figure).
When old, they have all shrunk to the same radius,
about that of the planet Jupiter.
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We have met them already in Section 9.2c.
For decades, there was a debate as to whether brown
dwarfs exist, but finally some were found in 1995.
We now know of about 1000, because of the advances
in astronomical imaging and in spectroscopy, not only
in the visible but also in the infrared.
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The coolest ones, of spectral type T, show methane and
water in their spectra, like giant planets but unlike
normal stars.
12.6 Brown Dwarfs
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It is difficult to tell the difference between a brown dwarf
and a small, cool, ordinary star, unless the brown dwarf is
exceptionally cool.
One way is to see whether an object has lithium in its
spectrum. Lithium is a very fragile element, and
undergoes fusion in ordinary stars, which converts it to
other things.
So if you detect lithium in the spectrum of a dim star, it is
probably a brown dwarf (which isn’t sustaining nuclear
fusion using protons) rather than a cool, ordinary dwarf
star of spectral class M or L, which are the coolest stars
on the main sequence (and thus have begun to sustain
their nuclear fusion).
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A complication is that very young M and L stars might not be
old enough to have burned all their lithium, leading to
potential confusion with brown dwarfs.
12.6 Brown Dwarfs


How do we tell the difference between brown dwarfs and giant
planets in cases where they are orbiting a more normal star?
Some astronomers would like to distinguish between them by
the way that they form: While planets form in disks of dust and
gas as the central star is born, brown dwarfs form like the
central star, out of the collapse of a cloud of gas and dust.



But we can’t see the history of an object when we look at it, so it is
hard to translate the distinction into something observable.
All of the proposed tests are difficult to make.
So, currently, for lack of definitive methods, the distinction is
usually made on the basis of mass: Any orbiting object with a
mass less than 13 times Jupiter’s is called a planet, while the
range 13 to 75 Jupiter masses corresponds to brown dwarfs.
(Objects less massive than 13 Jupiter masses not orbiting stars
are sometimes called “free-floating planets” since they are not
planets in the conventional sense of the word.)
12.6 Brown Dwarfs

The rationale for using 13 Jupiter masses as the dividing
line between planets and brown dwarfs is that above this
mass, fusion of deuterium occurs for a short time,
whereas below this mass, no fusion ever occurs.
 Thus, although brown dwarfs are not normal stars, they do
fuse nuclei for a short time, and hence aren’t completely
“failed stars” as many people call them.


Brown dwarfs are being increasingly studied, especially in
the infrared.
Hubble Space Telescope images show that one of the
nearby ones is double, with the components separated by
5 A.U.

By watching it over a few years, we should be able to
measure its orbit and derive the masses of the components.
12.7 The Solar Neutrino Experiment

Astronomers can apply the equations that govern matter and
energy in a star, and make a model of the star’s interior in a
computer.





Though the resulting model can look quite nice, nonetheless it
would be good to confirm it observationally.
However, the interiors of stars lie under opaque layers of gas.
Thus we cannot directly observe electromagnetic radiation from
stellar interiors.
Only neutrinos escape directly from stellar cores.
Neutrinos interact so weakly with matter that they are hardly
affected by the presence of the rest of the Sun’s mass.


Once formed, they zip right out into space, at (or almost at) the
speed of light.
Thus they reach us on Earth about 8 minutes after their birth.
12.7 The Solar Neutrino Experiment

Neutrinos should be produced in large quantities by the protonproton chain in the Sun, as a consequence of protons turning into
neutrons, positrons, and neutrinos; see the figure. (A positron is
an “antielectron,” an example of antimatter.

Whenever a
particle and its
antiparticle
meet, they
annihilate each
other.)
12.7a Initial Measurements

For over three decades, astrochemist Raymond Davis
has carried out an experiment to search for neutrinos
from the solar core, set up in consultation with the
theorist John Bahcall, whose calculations long drove
the theory.


One isotope of chlorine can, on rare occasions, interact
with one of the passing neutrinos from the Sun.


Davis set up a tank containing 400,000 liters of a
chlorine-containing chemical (see figure).
It turns into a radioactive form of argon, which Davis and his colleagues at
the University of Pennsylvania can detect.
He needs such a large tank because the interactions are so rare for a
given chlorine atom.

In fact, he detects fewer than 1 argon atom formed per day, despite the huge
size of the tank.
12.7a Initial Measurements

Over the years, Davis and his colleagues detected only about ⅓
the number of interactions predicted by theorists.





The latest thinking is that neutrinos actually change after they
are released.
According to a theoretical model, the neutrinos of the specific
type produced inside the Sun change, before they reach Earth,
into all three types (called “flavors”) of neutrinos that are
known.


Where is the problem?
Is it that astronomers don’t understand the temperature and
density inside the Sun well enough to make proper predictions?
Or is it that the physicists don’t completely understand what
happens to neutrinos after they are released?
But Davis’s experiment is sensitive only to the specific flavor
(“electron neutrinos”) released by the Sun.
Thus only ⅓ the original prediction is expected, and that is
what we detect.
12.7a Initial Measurements


The chlorine experiment, though it has run the longest by far, is no
longer the only way to detect solar neutrinos.
An experiment in Japan (see figure), led by Masatoshi Koshiba,
first set up to study protons and whether they decay, has used a
huge tank of purified water to verify that the number of neutrinos
is less than expected.

For their pioneering
work in the detection
of astrophysical
neutrinos, Davis and
Koshiba were
together awarded
half the 2002 Nobel
Prize in Physics.
12.7a Initial Measurements

Other sets of experiments in Italy and in Russia that use
gallium, an element that is much more sensitive to neutrinos
than chlorine, also show that up to half of the expected
neutrinos are missing.



Neutrinos were first thought of theoretically as particles that
have no rest mass— that is, particles that would have no mass
if they weren’t moving (“at rest”).



The chlorine was sensitive only to neutrinos of very high energy,
which come out of only a small fraction of the nuclear reactions in
the Sun and not out of the basic proton-proton chain.
Gallium is sensitive to a much wider range of interactions, including
the most basic ones.
The Japanese experiment has shown that neutrinos probably have
a tiny rest mass after all.
Theoretically, neutrinos can oscillate from one type to another
only if they have some rest mass.
So this result fits with the current ideas that we are seeing
neutrino oscillations from one type to another.
12.7b The Sudbury Neutrino Observatory


A U.S.–Canadian experiment in
Sudbury, Ontario, Canada, began
collecting data in 1999 (see figure).
It has even more sensitive detection
capability than earlier experiments.


It uses a large quantity of “heavy
water,” water whose molecules
contain deuterium instead of the
more normal hydrogen isotope.
This experiment, like those in Japan,
looks for light given off when
neutrinos hit the water.
12.7b The Sudbury Neutrino Observatory


The Sudbury Neutrino Observatory (SNO) has apparently
resolved the remaining mysteries about the missing solar
neutrinos.
They aren’t missing after all, since SNO has been able to
detect the correct rate in one of its configurations, which
is sensitive to all flavors of neutrinos.


In another mode, sensitive only to the electron neutrinos
that the Sun gives off, it confirms the deficit found by other
detectors.
So SNO has confirmed that most of the neutrinos change
in flavor en route from the Sun to the Earth.

The solar-neutrino problem was indeed in the neutrino
physics rather than in our understanding of the temperature
of the Sun’s core.
12.7b The Sudbury Neutrino Observatory


SNO will provide even more data over the coming years.
In particular, it should be able to determine the effect of
the Earth’s mass on neutrinos by comparing what it
detects during daylight hours with what it detects during
nighttime hours, when the neutrinos from the Sun have to
pass through the Earth to reach the detector.


The various neutrino experiments may cause a revolution in
fundamental ideas of physics.
Similarly, there are also important repercussions for
physics from a set of astronomical observations we will
discuss later (Section 18.5), showing that measurements
of distant supernovae (exploding stars) indicate that the
expansion of the Universe is accelerating.

The relation between physics and astronomy is close,
though the point of view that physicists and astronomers
have in tackling problems may be different.
12.8 The End States of Stars


The next two chapters discuss the various end states of stellar
evolution.
The mass of an isolated, single star determines its fate and we provide a
figure here as a type of “coming attractions” (see figure).

This brief section and
diagram thus set the
stage for what is
coming next, which is
so interesting that this
brief introduction
leads to two separate
chapters.
12.8 The End States of Stars


In Chapter 13, we discuss low-mass stars, like the Sun,
which wind up as white dwarfs.
We also discuss the more massive stars, which use up
their nuclear fuel much more quickly.


These high-mass stars are eventually blown to smithereens
in supernova explosions, becoming neutron stars or even
black holes!
Chapter 14 is devoted entirely to exotic black holes and
their properties.