How the Stars Shine:

1
Chapter 12

• How the Stars Shine
• Cosmic Furnaces

2
Introduction

• 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.
• 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.
• However, we must take into account the different
  masses of stars some aspects of their evolution
  depend critically on mass.

3
Introduction

• 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.

4
Introduction

• 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.
• 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.
• These astronomical results therefore have added
  important knowledge about fundamental physics in
  addition to our understanding of the stars.

5
12.1 Starbirth

• The birth of a star begins with a nebulaa large
  region of gas and dust (see figures).
• The dust (tiny solid particles) may have escaped
  from the outer atmospheres of giant stars.

6
12.1 Starbirth

• 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.
• We discuss the infrared observations largely
  here, including the new capabilities of NASAs
  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).
• 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.

7
12.1a Collapse of a Cloud

• Consider a region that reaches a higher density
  than its surroundings, perhaps from a random
  fluctuation in density orin a leading theory of
  why galaxies have spiral armsbecause 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.
• 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.

8
12.1a Collapse of a Cloud

• 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.)

9
12.1a Collapse of a Cloud

• The rising internal pressure, which is highest in
  the protostars 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).

10
12.1a Collapse of a Cloud

• 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
  acrossthat 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.
• 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.

11
12.1a Collapse of a Cloud

• 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-main-sequence star, its internal pressure
  decreases, and it gradually contracts.
• 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.

• 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.
• As we shall discuss later, this mechanical
  balance is the key to understanding stable stars
  see Figure at left.

12
12.1a Collapse of a Cloud

• 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.
• 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.

13
12.1a Collapse of a Cloud

• 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.
• 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).

14
12.1a Collapse of a Cloud

• 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.
• We see this evaporation taking place at the
  tops of the Eagle Nebulas pillars.
• The stars are destroying their birthplaces as
  they become independent and more visible from
  afar.

15
12.1b The Birth Cries of Stars

• 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.

16
12.1b The Birth Cries of Stars

• 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).
• The ejection of these spinning clumps helps slow
  the stars 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.
• 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?

17
12.1b The Birth Cries of Stars

• 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.

• Since the pre-main-sequence star is hidden in
  visible light by a dusty cocoon of gas, infrared
  observations of Herbig-
• Haro objects most clearly reveal what is going
  on (see figure).

18
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 premain-sequence star, the gas
  squirts outward along the pre-main-sequence
  stars axis of rotation at speeds of perhaps 1
  million km /hr.
• HH1 and HH2 (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.

• These objects are about 1500 light-years from us,
  in a star-forming 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.

19
12.1b The Birth Cries of Stars

• Several classes of stars that vary erratically in
  brightness have been found.
• 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 havent 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.
• 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.

20
12.2 Where Stars Get Their Energy

• 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 Suns light and heat,
  date back billions of years.
• Some other source of energy must hold the Sun and
  other stars up against their own gravitational
  pull.

21
12.2 Where Stars Get Their Energy

• 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.
• 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 stars luminosity and temperature change little
  while it is on the main sequence nuclear
  reactions provide the stability.

22
12.2 Where Stars Get Their Energy

• The energy makes the particles in the star move
  around rapidly.
• 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.
• This outward force balances gravitys inward pull
  on the pocket (hydrostatic equilibrium, which
  we illustrated in the figure).

23
12.2 Where Stars Get Their Energy

• 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.)

• 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.
• A small amount of the mass, m, disappears in
  the process 0.007 (0.7 per cent) of the mass of
  the four protons.

24
12.2 Where Stars Get Their Energy

• The mass difference does not really disappear,
  but rather is converted into energy, E, according
  to Albert Einsteins famous formula
• E mc2,
• where c is the speed of light.
• 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.
• This energy is known as the binding energy of
  the nucleus, here specifically that of helium.
• 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.

25
12.2 Where Stars Get Their Energy

• 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.

26
12.3 Atoms and Nuclei

• 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.

27
12.3a Subatomic Particles

• 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.
• 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).

28
12.3a Subatomic Particles

• 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.
• 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.
• The element with one proton is hydrogen, that
  with two protons is helium, that with three
  protons is lithium, and so on.

29
12.3b Isotopes

• 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.
• 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.
• Another isotope of hydrogen called tritium has
  one proton and two neutrons.

30
12.3b Isotopes

• 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).
• 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?92146 neutrons.
• Each element has only certain isotopes.
• For example, most of the naturally occurring
  helium is in the form 2He4, with a much lesser
  amount as 2He3.

31
12.3c Radioactivity and Neutrinos

• 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.
• 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.)

32
12.3c Radioactivity and Neutrinos

• 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.
• 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.
• 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.

33
12.3c Radioactivity and Neutrinos

• The elusiveness of the neutrino not only makes it
  a valuable messengerindeed, the only possible
  direct messengercarrying 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 ? the expected number of
  neutrinos, as we shall soon see.

34
12.4 Stars Shining Brightly

• 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.
• 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.
• 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.

35
12.4 Stars Shining Brightly

• However, at the high temperatures (millions of
  kelvins) typical of a stellar interior, some
  nuclei occasionally have enough energy to
  overcome this electrical repulsion.
• 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.
• 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 gravitys inward pull.

36
12.4 Stars Shining Brightly

• In the center of a star, the fusion process is
  self-regulating.
• The star finds a balance between thermal pressure
  pushing out and gravity pushing in.
• 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 Earths oceans.

37
12.4 Stars Shining Brightly

• The greater a stars 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.
• 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.
• The next two chapters examine the ultimate fates
  of stars, with the fates differing depending on
  the masses of the stars.

38
12.5 Why Stars Shine

• 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.

39
12.5 Why Stars Shine

• 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 isnt the
  main point.)

40
12.5 Why Stars Shine

• 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 Emc2.
• According to Einsteins special theory of
  relativity, mass and energy are equivalent and
  interchangeable, linked by this equation.
• For stellar interiors significantly hotter than
  that of the Sun, the carbon-nitrogen oxygen (CNO)
  cycle dominates.
• 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.
• Thus, as much carbon remains at the end as there
  was at the beginning, and the carbon can start
  the cycle again.

41
12.5 Why Stars Shine

• 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
  Emc2.
• Main-sequence stars more massive than about 1.1
  times the Sun are dominated by the CNO cycle.
• 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.
• Since three helium nuclei (2He4) go into making a
  single carbon nucleus (6C12), the procedure is
  known as the triple-alpha process.

42
12.5 Why Stars Shine

• A series of other processes can build still
  heavier elements inside very massive stars.
• These processes, and other element-building
  methods, are called nucleosynthesis
  (newclee-oh-sintha-sis).
• 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.

43
12.6 Brown Dwarfs

• When a pre-main-sequence star has at least 7.5
  per cent of the Suns 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
  Suns 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
  doesnt last long because there is so little
  deuterium in the Universe relative to ordinary
  hydrogen.)

44
12.6 Brown Dwarfs

• 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 didnt 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.
• 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.
• The coolest ones, of spectral type T, show
  methane and water in their spectra, like giant
  planets but unlike normal stars.

45
12.6 Brown Dwarfs

• 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 isnt
  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).
• 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.

46
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 cant 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 Jupiters 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.)

47
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
  arent 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.

48
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 stars 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 Suns 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.

49
12.7 The Solar Neutrino Experiment

• Neutrinos should be produced in large quantities
  by the proton-proton 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.)

50
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.
• Davis set up a tank containing 400,000 liters of
  a chlorine-containing chemical (see figure).
• One isotope of chlorine can, on rare occasions,
  interact with one of the passing neutrinos from
  the Sun.

• 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.

51
12.7a Initial Measurements

• Over the years, Davis and his colleagues detected
  only about ? the number of interactions predicted
  by theorists.
• Where is the problem?
• Is it that astronomers dont understand the
  temperature and density inside the Sun well
  enough to make proper predictions?
• Or is it that the physicists dont completely
  understand what happens to neutrinos after they
  are released?
• 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.
• But Daviss experiment is sensitive only to the
  specific flavor (electron neutrinos) released
  by the Sun.
• Thus only ? the original prediction is expected,
  and that is what we detect.

52
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.

53
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.
• 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.
• Neutrinos were first thought of theoretically as
  particles that have no rest mass that is,
  particles that would have no mass if they werent
  moving (at rest).
• 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.

54
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.

55
12.7b The Sudbury Neutrino Observatory

• The Sudbury Neutrino Observatory (SNO) has
  apparently resolved the remaining mysteries about
  the missing solar neutrinos.
• They arent 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 Suns core.

56
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 Earths 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.

57
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.

58
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.