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The Lives and Deaths of
Low-Mass Stars
Courtesy of NASA, ESA, HEIC, and the Hubble Heritage Team (STScI/AURA)
Chapter 14
14-1 Brown Dwarfs
1. The maximum mass that a star can have is about 100
solar masses.
– Any more and the internal pressure and huge
amounts of energy generated in the core would
overwhelm gravity and blow the star apart.
2. The minimum mass that a star can have is about
0.08 solar masses.
– Any less and its core cannot become hot and
dense enough for nuclear reactions to begin.
3. A brown dwarf is a star-like
object whose mass is too
small to sustain nuclear
fusion.
– Probable limits of mass
for brown dwarfs are from
0.002 to 0.08 solar
masses. (Jupiter’s mass is
about 0.001 solar masses.)
– After a brief initial phase
where some deuterium is
burned, brown dwarfs
continue to cool.
Figure 14.01a
4. A brown dwarf
(GL229B) was first
seen in 1994. Its
mass is between 20
and 70 times that of
Jupiter.
Figure 14.01c: The brown dwarf 2M1207a and its companion planet
2M1207b
Courtesy of Paranal Observatory, ESO
5. Most known brown dwarfs are isolated. Their
formation is either similar to that of a normal
isolated star or they are the byproduct of a chaotic
star formation process during which they are
ejected from a new-born stellar group.
6. A brown dwarf is not a star because it cannot
sustain nuclear reactions in its core.
14-2 Stellar Maturity
Stellar Nuclear Fusion
1. In stars of low mass (less than 1.5 solar masses),
the predominant energy-generating series of
reactions is the proton-proton (p-p) chain.
2. In stars of mass greater than 1.5 solar masses,
the core temperatures are greater and thus a
different chain of nuclear reactions (the CNO
cycle) dominates.
This series of reactions involves
– hydrogen,
– carbon,
– nitrogen,
– and oxygen.
3. The carbon or CNO cycle is a series of nuclear
reactions that results in the fusion of four
hydrogen nuclei into a helium nucleus (just like the
p-p chain), using carbon-12 as a catalyst in the
process.
Figure 14.03: The CNO cycle changes four hydrogen nuclei into one helium
nucleus, with an attendant release of energy.
The Stellar Thermostat
1. The core of a main sequence star has a
regulating mechanism (just like a thermostat)
that controls the rate of consumption of
hydrogen fuel.
2. The overall effect of the mechanism is that the
star remains in hydrostatic equilibrium.
Main Sequence Life
1. In the core of a main sequence star, the following
sequence of events occurs:
Fig. 14-4
2. As a star ages on the
main sequence,
– it becomes more
luminous,
– its radius increases,
– its surface (effective)
temperature
decreases.
The star’s position on
the H-R diagram moves
to the up and right of
the main sequence.
Fig. 14-5
3. Massive stars have a
greater fusion rate and
thus are more
luminous.
– As a result, they can
use up the hydrogen
in their cores in a
few million years.
– The least massive
stars can sustain
hydrogen fusion for
hundreds of billions
of years.
Figure 14.06a: An H-R diagram of the
Pleiades reveals that its most massive
stars have started leaving the main
sequence.
Figure 14.06b: The cluster M11 is older than the Pleiades, and stars are
turning off the main sequence at a lower point.
4. Observational evidence for the shorter lifetimes
of massive stars comes from galactic clusters.
Assuming stars in a cluster formed at about the
same time, the older the cluster, the farther its
most massive stars have moved to the right from
the zero-age main sequence.
5. The main sequence lifetime of the stars at the
turnoff point (the point on the H-R diagram of a
cluster of stars where the stars are just leaving
the main sequence) is equal to the cluster’s age.
Tools of Astronomy: Lifetimes on the Main
Sequence
1. The amount of time a star spends on the main
sequence depends on the amount of hydrogen
in its core and its rate of hydrogen
consumption.
2. The lifetime (t ) of a star on the main sequence is
given by the expression:
t = tSun / M 2.5
where M is the star’s mass in solar units and
TSun is the Sun’s main sequence lifetime
(about 10 billion years).
14-3 Star Death
1. Until their lives end on the main sequence, the main
difference between the evolution of stars of various
masses is the amount of time they spend as
protostars and main sequence stars.
2. After this point, the star’s mass determines which of
several different paths its life will take.
3. Stars can be grouped by mass as
– very low mass (< 0.4 M),
– moderately low mass (0.4–4 M),
– moderately massive (4–8 M),
– and very massive (> 8 M).
14-4 Very Low Mass Stars (< 0.4 M)
1. In stars with a mass of less than about 0.4 solar
masses, convection occurs throughout most or all
of the star’s volume.
Hydrogen from throughout the star is cycled
through the core, and the entire star runs low on
hydrogen at the same time.
Figure 14.07: Convection in
stars of various masses
2. As the rate of fusion decreases
in the core, the star contracts
and heats up. The position of
the star on the H-R diagram
moves toward the lower left,
and the star becomes a white
dwarf.
3. A very low mass star will take
more than 20 billion years to
completely burn its hydrogen.
4. Since the lifetime of a very low
mass star is more than the age
of the universe, white dwarfs
currently observed must have
originated in a different manner.
Figure 14.08: H-R diagram
of low-mass stars
14-5 Beyond the Very Low Mass Stars:
The Red Giant Stage
1. About 90% of the stars in the sky are on the main
sequence.
This implies that a typical star spends 90% of its
luminous lifetime on the main sequence.
2. For most of these stars the next step after
leaving the main sequence is to become red
giants.
3. A typical star begins to contract once its core is
depleted of hydrogen. This gravitational
contraction causes the star to heat up.
4. Energy from the
contraction of the core
then heats a shell
surrounding the core to
temperatures that
permit fusion of
hydrogen to begin.
5. These two sources of
energy cause the outer
portions of the star to
expand and cool.
Fig. 14-9
6. The position of the star on the H-R diagram
moves to the right and upward (due to
increasing luminosity).
7. A red giant can have a lower surface
temperature (less radiation per square meter)
but a higher luminosity because its diameter
will have expanded by a large factor (200 times
for a “typical” red giant).
14-69 Moderately Low Mass Stars (0.4 – 4 M)
Electron Degeneracy and the Helium Flash
1. The core of a red giant consists of helium nuclei
intermingled with electrons, a mixture that has
properties similar to a regular gas.
2. The core of a red giant will not continue to contract
indefinitely because of electron degeneracy.
3. Electron degeneracy is the state of a gas in
which its electrons are packed as densely as
nature permits.
– The temperature of such a high-density gas is not
dependent on the pressure as it is in a “normal” gas.
4. In the case of the degenerate core of a red giant,
the more massive the core is, the smaller is its
radius.
5. As the red giant evolves and hydrogen burning
takes place in its outer layers, the helium “ashes”
are dumped back onto the degenerate core,
raising the temperature of the core.
Figure 14.11a: For normal, main sequence stars, the more massive the star
is, the larger is its radius.
Figure 14.11b: For the degenerate core of a red giant, the more massive
the core is, the smaller is its radius.
6. When the core temperature reaches 100 million
Kelvin, helium nuclei begin to combine, forming
carbon.
Figure 14.12: The steps that occur as helium fuses into carbon in the core
of a red giant
7. Because of degeneracy, the core cannot cool by
expanding and thus helium fusion reactions occur
faster and faster.
– Helium flash is the process of runaway helium
fusion reactions that occurs during the evolution
of a red giant.
8. The end result of the helium flash is that
– the core heats up,
– the degeneracy is destroyed,
– the stellar thermostat is restored,
– and some mass from the star’s surface may
be lost.
9. Following the helium flash, the star contracts
slightly and its surface temperature increases.
10. The by-products of helium fusion reactions are
carbon and oxygen.
When a red giant reaches the stage where it has
a carbon core,
– the heat from the shrinking core ignites
helium fusion in a shell around it,
– while hydrogen is fusing in a second shell
beyond the first.
11. The star slowly enters a new red giant phase
and its position on the H-R diagram moves to the
right and upward.
Figure 14.13: Shells of hydrogen and helium fusion in a red giant's shrinking core
Figure 14.14: When the Sun becomes a red giant, it will expand until its
surface is somewhere between Earth and Mars.
12. Stars more massive than 2 solar masses do not
experience a helium flash.
– Their cores make a smooth transition to
helium burning without becoming
degenerate.
– When their helium supply is exhausted,
however, their internal structure is similar to
that of Sun-like stars.
Stellar Pulsations
1. During its evolution, a star continuously tries to
remain in equilibrium. The changes occurring in
its core are often periodic.
2. Two types of pulsating stars are especially
important to astronomers because they can be
used as distance indicators:
– the Cepheid variables (with periods between 1
and 100 days)
– the RR Lyrae variables (with periods shorter
than one day)
3. The oscillations of the star’s outer layers
result from the ionization of helium in the
outer layers when it gets compressed.
– The resulting gas is opaque, trapping
heat from inside the star.
– This results in an increase of the
pressure under the layers, an expansion
and subsequent cooling of the layers.
– The layers then become transparent, the
trapped energy is released, the layers fall
inward and the cycle repeats itself.
4. The conditions necessary for such pulsations
occur in a narrow strip on the H-r diagram
called the instability strip.
5. Cepheid variables result from massive stars in
the instability strip, while RR Lyrae variables
result from low mass stars.
6. In some cases stellar pulsations can be so
large that the star loses its outer layers to
outer space.
Figure 14.15: A section of the instability strip above the main sequence.
Mass Loss from Red Giants
1. Observations show that other main sequence stars
shed material just like our Sun does as the solar
wind.
2. The solar wind carries away about 10–14 of the
Sun’s mass each year.
– Over the course of 10 billion years, the Sun will
lose only 0.01% of its mass this way.
3. In red giant stars it is thought that core instabilities
and pulsations are responsible for the large mass
loss.
– A typical red giant loses 10–7 solar masses a
year.
Courtesy of NASA/ESA/JPL-Caltech/J. Hora (CfA) & C.R. O'Dell (Vanderbilt)
Planetary Nebulae
1. A planetary nebula is
a spherical shell of
gas that is expelled
by a low mass red
giant near the end of
its life.
Figure 14.17b: The Helix nebula (NGC 7293)
2. Planetary nebulae have nothing to do with
planets. They were named so decades ago
because their colors resemble the colors of
Neptune and Uranus when seen through a small
telescope.
3. The material in the shell glows because UV
radiation from the central hot star causes it to
fluoresce.
4. Pulsations in the core of the red giant and/or
stellar winds emitted from the dying star (and
which occur in definite stages) are thought to
cause planetary nebulae.
5. Not all planetary
nebulae appear as
rings.
– Some show hints of
spherical structure
while others are
bipolar.
Figure 14.19a: The Dumbbell nebula (M27)
Courtesy of NASA/ESA/JPL-Caltech/J. Hora (CfA) & C.R. O'Dell (Vanderbilt)
– Their shapes depend
on how the outflows
interact with each
other and with the
star’s magnetic field.
Figure 14.DP03: Hubble 5 is a butterfly or
bipolar nebula
Courtesy of B. Balick (U. Washington) et al., WFPC2, NASA, HST, NASA
Figure 14.DP04: NGC7009
Courtesy of B. Balick (.U. Washington), et.al, WFPC2, HST, NASA
Figure 14.21a: The straight features in Red Rectangle imply episodes of
mass ejection from the central object.
Courtesy of NASA, ESA, Hans Van Winckel (Catholic University of
Leuven, Belgium), and Martin Cohen (University of California, Berkeley)
6. Planetary nebulae do not last long.
Their material quickly dissipates into space and
becomes part of the interstellar medium.
7. The proto-planetary nebula phase lasts between
a few hundred to 1000 years.
(This phase corresponds to the transition between the
last stages of a red giant star’s life and the late stages
of a planetary nebula after the star ejected most of its
mass.)
Figure 14.24: An H-R diagram expanded to the left—to higher
temperatures—to include stars passing through the planetary nebula stage.
8. Observations show that proto-planetary nebulae
are huge organic chemical factories.
9. The core left behind after the nebula dissipates
becomes a white dwarf.
14-7 White Dwarfs
1. Sirius B was the first white dwarf observed (in
1862); it has about the same mass as Sirius A, but is
small and hard to see.
2. Astronomers estimate that 10% of all stars are white
dwarfs.
3. White dwarfs are the cores of red giants that remain
after the outer parts of the original stars have blown
away.
Figure 14.25b: Cluster of stars with 7 white dwarfs
Kitt Peak National Observatory 0.9-meter telescope, National Optical Astronomy Observatories; Courtesy M. Bolte (University of California, Santa
Cruz) and Harvey Richer (University of British Columbia, Vancouver, Canada) and NASA
The Chandrasekhar Limit
1. Even though electron degeneracy supports the
white dwarf against collapsing completely, there
is a limit to the amount of pressure degenerate
electrons can withstand.
2. This limit to the mass of a white dwarf above
which it cannot be supported by electron
degeneracy and cannot exist as a white dwarf is
known as the Chandrasekhar limit.
– It is about 1.4 solar masses.
3. Main sequence stars with masses up to 4 M can
end up as white dwarfs only if they lose mass
during the red giant and planetary nebula phases.
Characteristics of White Dwarfs
1. White dwarfs that formed from low mass stars
are composed of helium and carbon.
White dwarfs that formed from more massive
stars contain nuclei of oxygen, neon, sodium,
and even iron.
2. White dwarfs have observed surface
temperatures between 4,000 K and 85,000 K.
Their masses range from perhaps 0.02 solar
masses up to 1.4 solar masses.
3. A typical white dwarf will
have
– 0.8 solar masses,
– a diameter of 10,000
km (3/4 of Earth’s),
– and a density of 106
g/cm3.
A teaspoon of white
dwarf material would
weigh two tons.
Figure 14.26: When the Sun becomes a white
dwarf, it will be slightly larger than today’s Earth.
4. To get to the white dwarf stage, a low-mass star
will have gone through these stages:
protostar,
main sequence star,
red giant,
planetary nebula.
5. White dwarfs do not produce
energy. They radiate away
their leftover energy and
simply fade away and
become black dwarfs.
6. Black dwarf is the theorized
final state of a star with a
main sequence mass less
than about 8 solar masses,
in which all of its energy
sources have been depleted
so that it emits no radiation.
Figure 14.27: Summary of stellar evolution
Novae
1. A binary system of a white dwarf and a newly
formed red giant will result in the formation of an
accretion disk around the white dwarf. The
material in the disk comes from the red giant and
is mostly hydrogen.
2. An accretion disk is a rotating disk of gas orbiting
a star, formed by material falling toward the star.
Figure 14.28
3. The hydrogen builds up on the surface of the
white dwarf, becomes denser and hotter and can
ignite in an explosive fusion reaction when the
temperature reaches 10 million K.
4. This surface explosion blows off the outer layers
of the white dwarf.
Though this shell contains a tiny amount of mass
(0.0001 solar masses) it can cause the white
dwarf to brighten by 10 to 20 magnitudes (10,000
to 100 million times brighter) in a few days.
5. A nova is a star that suddenly and temporarily
brightens, thought to be due to new material
being deposited on the surface of a white dwarf.
6. Because so little mass is blown off during a
nova, the explosion does not disrupt the binary
system.
Ignition of the infalling hydrogen can recur with
periods ranging from months to thousands of
years.
14-8 Type I Supernovae
1. If accretion brings the mass of a white dwarf
above the Chandrasekhar limit, electron
degeneracy can no longer support the star, and
it collapses.
– The collapse raises the core temperature,
– runaway carbon fusion begins,
– which ultimately leads to the star exploding
completely.
2. Such an exploding white dwarf is called a
supernova.
3 . While a nova may reach an absolute magnitude of
–8 (about 100,000 Suns), a supernova attains a
magnitude of –19 (10 billion Suns).
4. There are two types of supernovae:
(a) Type I
– their spectrum has no hydrogen lines.
(b) Type II
– their spectrum contains prominent hydrogen
lines
– they originate from the explosion of a single
star.
5. Type I supernovae are divided into three
subclasses:
(a) Type Ib, and Ic are caused by massive stars
that have lost different proportions of their outer
layers before exploding.
(b) Type Ia result from white dwarfs.
6. A Type Ia supernova
– reaches maximum brightness in a few days,
– fades quickly for about a month,
– and then declines in brightness more
gradually until it dissipates in about a year.
7. Models indicate that the energy of a Type Ia
supernova (following the explosion) comes from
radioactive decay of nuclei produced in the
explosion.
8. The light curves and spectra of all Type Ia
supernovae are very similar.
This similarity is a powerful tool that astronomers
use to trace the history of the expansion of the
universe.
Figure 14.30: Light curve of
Type Ia supernova
Figure 14.31: (left) X-ray data of the supernova remnant DEM L71. (right)
The blast wave at optical wavelengths.
X-ray: NASA/CXC/Rutgers/J.Hughes et al.; Optical: Rutgers Fabry-Perot
Figure 14.32: The Vela nebula is a Type II supernova remnant
© Royal Observatory, Edinburgh/AATB/SPL/Photo Researchers, Inc.