Lec 6
II-8 Stellar Evolution
(Main Ref.: Lecture notes; Parts of FK
Sec. 16-1,2, Ch 18, 19, 20; CD photos
shown in class)
II-8a. Introduction (Main Ref.:
Lecture notes; FK p.505, p. 533, p. 563,
Sec. 18-1)
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Understanding how stars evolve requires both
observation and ideas from physics
• Because stars shine by thermonuclear
reactions, they have a finite life span
• The theory of stellar evolution describes
how stars form and change during that life
span
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II-8b. Early Stages (Star Formation)
and Middle Age (Main Sequence)
(Main Ref.: Lecture notes; FK, Sec. 16-1,2, 18-2 through 8, 19-1)
(i) Interstellar Nebulae (Main Ref.: Lecture FK Sec. 182,3;CD photos shown in class)
Interstellar gas and dust pervade the Galaxy
•The interstellar matter (ISM) filling the space between stars is not
homogeneous nor uniform, but clumpy – there are denser
regions like clouds. These denser regions are generally called
interstellar `Nebulae’. (Note: there are other kinds of nebulae,
too, e.g., planetary nebulae, matter surrounding supernovae,
etc., but we will come back to them later.)
• There are three kinds of interstellar nebulae - dark nebulae,
emission nebulae, and reflection nebulae. We will briefly
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summarize these nebulae.
• Dark Nebulae - Black:
• Dark nebulae are cool gas and dust so dense that they are
opaque. They appear as dark blots against a background
of distant stars when they are in our line of sight
•  stars are born there!
• Examples:, Barnard objects, Bok globules, Horsehead
nebula in Orion, Barnard 86.
See CD pictures in class, and Fig. II-42, 43, 45, 46, 47, 48.
• Why dark? Because dust grains and cool gas are dense
enough to block the way of light from background stars,
emission nebulae, HII regions, etc. (See class notes for
how it works.)
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Emission Nebulae – Bright Red:
• Emission nebulae, or H II regions, are glowing, ionized hot
clouds of gas.
Emission nebulae are powered by ultraviolet light that they
absorb from nearby hot stars
Example: Orion Nebula, Rosette Nebula
See CD pictures in class, and Fig. II-42, 43, 48.
Why Bright Red? Because near-by bight, hot young stars,
e.g., O, B stars just born, illuminate a cloud of hot gas by
UV light and heat up the gas. Most abundant element of the
gas (~ 74%) is Hydrogen (H). When the gas heats up, the
neutral H atom (= HI) is ionized (electron escapes) and
become singly charged H ion, H+ (= proton p = HII).
When H ion becomes H atom again by recombination, 5
photons are emitted in a cascade. The recombined electrons are now
bound to proton, and jump down from higher energy levels to lower
ones, emitting Balmer emission lines, which are in the visible
range. The most prominent line is H, which comes in red, and so
Red Color! (See class notes for further details.)
• Reflection Nebulae - Blue:
• Dense, cool cloud of gas and dust grains, NOT in our line
of sight, reflects light from near by stars, HII regions, etc.
Example: NGC 2023 in Orion Nebula, Pleiades.
See CD pictures in class, and Fig. II-43, 44, 48.
Why blue? Because the dust grains and gas in the dense
cloud are too cool to emit light by itself (unlike emission
nebulae), but can reflect light from near-by stars and HII
regions. We can see the reflected light if the cloud is not
in our line of sight. Looks blue because more blue is
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reflected – the same as why the sky is blue.
Note: If the cloud is in our line of sight and if it is dense
enough, the cloud will appear as a dark nebula!
(See class notes for how it works.)
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Fig II-42: Orion Nebula
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EX 29
Fig II-43: Orion Nebula
Fig II-44: NGC 6726-27-29
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EX 30
Fig II-45: a Dark Nebula
Fig II-46: Bok Globules9
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Fig II-47: the Milky Way
Fig II-48: the Trifid Nebula
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(ii) Star Formation – Protostars and Birth
(Main Ref.: Lecture notes; FK Sec. 18-3 to 8; CD photos shown in
class)
• As already noted in Section (i), the interstellar matter (ISM) is
not uniform, but clumpy. New stars are formed in these
clumpy, cool, dense clouds called `dark nebulae’ in or near
molecular clouds (cool clouds with CO and H2 molecules).
Bursts of protostar formation takes place when these dense
regions are hit by high speed (`supersonic’, meaning speed
faster than the sound speed) winds from a near-by supernova
explosion or UV light and winds from near-by hot O and B stars
 shocks  protostar formation.
• Physically, a dense, cool cloud of gas starts to contract due to
self-gravity when a given amount of mass (hundreds of solar
mass) gets smaller in size than a critical radius, called `Jean’s
Radius’, RJ. During contraction the cloud fragments into
smaller clouds  Birth of protostars!
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• When a protostar contracts, gravitational energy is released.
According to the Virial Theorem (to be discussed in an upper
division astrophysics course, but not in this course), half of the
released gravitational energy is radiated away from the surface
as photons, while the other half heats up the star. So, as the
contraction proceeds, the protostar gets hotter and smaller.
• The protostar gets hot and luminous enough to shine when
the temperature gets to about 3000 K (so it should look red).
• This newborn protostar is surrounded and protected by a
dense cocoon of grains and gas which absorbs visible light
from the star, and so a protostar is hard to see by visible
light, but it can be seen by infrared (IR) light, because IR is
more transparent through the cocoon.
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• Since all protostars of all mass range are born at roughly
~ 3000 K, with different luminosities, they all lie along
the Hayashi Track on the HR Diagram. See Fig. II-56.
(See class notes for the details.)
• As the protostar keeps contracting and heats up, it
moves on the H-R Diagram from the Hayashi Track to
the main sequence. By the time it reaches the main
sequence the central temperature gets to > ~ 107 K, high
enough for the H burning to start. Then, the nuclear
energy can supply the source of energy lost by photons
from the stellar surface, and hence, the contraction stops
 Birth of a Star!
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EX 32
•Protostars form in cold, dark
nebulae
•As a protostar grows by
the gravitational accretion
of gases, Kelvin-Helmholtz
contraction causes it to
heat and begin glowing
•Star formation begins in
dense, cold nebulae,
where gravitational
attraction causes a clump
of material to condense
into a protostar
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Fig II-49: Newborn Stars in Orion Nebula
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Fig II-50: M16, a Star Cluster with Star Forming Regions
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Fig II-51: How O and B Stars Trigger
Star Formation
Star-forming regions appear when a giant
molecular cloud is compressed.
This can be caused by the cloud’s passage through
one of the spiral arms of our Galaxy, by a supernova
explosion, or by other mechanisms
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During the birth process, stars both gain
and lose mass
In the final stages of pre–main-sequence contraction, when
thermonuclear reactions are about to begin in its core, a
protostar may eject large amounts of gas into space
Low-mass stars that vigorously eject gas are called T Tauri
stars
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Fig II-52: Mass Loss fromYoung, Massive Stars
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O and B Stars and Their Relation to H II Regions
Fig II-53: Mapping Molecular Clouds
•The most massive protostars
to form out of a dark nebula
rapidly become main
sequence O and B stars
•They emit strong ultraviolet
radiation that ionizes
hydrogen in the surrounding
cloud, thus creating the
reddish emission nebulae
called H II regions
•Ultraviolet radiation and
stellar winds from the O and B
stars at the core of an H II
region create shock waves
that move outward through
the gas cloud, compressing
the gas and triggering the
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formation of more protostars
Supernovae compress the interstellar medium
and can trigger star birth
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Fig II-54: SN 0103-72.6
Fig II-55: An Association
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• Path of protostars from the Hayashi Track to Main
Sequence (MS):
Figure II-56 Hayashi Track
To understand how a
protostar moves from the
Hayashi Track to the main
sequence, it is convenient to
use the Stefan-Boltzman Law:
L = 4  R2  T4. Eqn(16)
High mass star
(e.g., 9M☉ star)
~ horizontal toward left –
R decreases while Ts (surface
temperature) increases in such
a way as to keep L ~ constant
(see Eqn(16)).
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Low mass star (e.g., Sun)
~
downward first – while R decreases, Ts increase is so small that L decreases (see
Eqn(16), but then, moves toward left before going down again slightly – finally
increase in Ts catches up with R decrease to keep L ~ constant, and
then reverses the trend again in the end, before reaching MS.
See Fig. II-56 and class notes for details.
Protostar Lifetime:
•
•
•
Short compared with the MS life time (age).
Shorter for more massive stars.
Why? More massive stars involve larger gravity  contract faster 
reach high enough temperature for H-burning faster!
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M = 1 M☉.  t ~ 2 x 107 years,
compare with t(MS) ~ 1010 years
M = 15 M☉  t ~ 105 years,
compare with t(MS) ~ 107 years
where t(MS) = main sequence lifetime.
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(iii) Middle Age – Main Sequence (MS)
Stars (Main Ref.: Lecture notes; FK Sec. 16-1,2, 18-4,
19-1)
• Stable and long, because:
(i) Nuclear energy source, H-burning, is stable and
lasts long 
no gravitational contraction 
mechanically balanced 
Gravity supported by gas pressure (ideal gas) .
Note: Nuclear energy source ( = H-burning) stable
and lasts long.
(ii) Energy balance and transport:
Nuclear energy released in central core = Radiation energy
lost from the surface by photons, through transport from
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center to surface.
★Mass:
Minimum Mmin ~ 0.08 M☉: Why? Because if M < Mmin ,
T too low (< 107 K) for H-burning  no nuclear energy source
 star cannot shine by itself - cannot be born as a star!
Maximum Mmax ~ 200 M☉: Why? Because if M > Mmax 
raiation pressure > gravity  hydrostatic equilibrium (balance) lost
 star cannot be held together and exist as a stable star.
★ Zero Main Sequence (MS) Star: = MS star at t = 0 i.e., when star was born = when H-burnng starts.
Composition at t = 0: H ~74%; He ~ 25%; others (Z > 2) ~ 1%,
for stars like the sun.
★ Non-Zero Main Sequence (MS) Star: = MS star at t > 0.
e.g., Our Sun: M = 1 M☉; Age t = 4.6 x 109 year.
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Composition of non-zero MS stars:
e.g., Center of our Sun: H ~30%; He ~ 70%; others (Z > 2) ~ 1%
See Fig. II-57 (below) and class notes
H mass fraction (%)
He mass fraction (%)
Fig. II-57: Internal composition of the sun
Note: during t = 0 to tend, R increases by ~ 6%, L increases by ~ 40%,
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and Ts increases by 300K.
Lifetime (duration) for MS ~ Age of Star: Long, but shorter for
more massive stars. Why? More massive MS stars are hotter 
H-burning proceeds more
quickly, energy radiated more
quickly, fuel (H) exhausted
more quickly, etc.
(see class notes & HR Diagram).
Exmples:
M = 25 M☉, tend ~ 4 x 106 years;
M = 15 M☉, tend ~ 1.5 x 107 years;
M = 3 M☉, tend ~ 8 x 108 years
M= 1 M☉, tend ~ 1.2 x 1010 years
M = 0.5 M☉, tend ~ 7 x 1011 years.
See class notes and Fig. II-58
For detailed explanation.
FK
Fig. II-58: Stellar evolution of
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MS and post-MS stars
H-R Diagram: .MS star moves up and to the right or left somewhat.
Hence, MS stars lie on the strip from t = 0 to tend = end of MS.
Study Fig. II-58 and class notes.
A star’s lifetime on the main sequence is
proportional to its mass divided by its luminosity
Table II-5: Main Sequence Lifetime
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The duration of a star’s main sequence lifetime depends on the rate at
which the hydrogen is consumed in the core and energy lost by
radiation from the surface
•The more massive a star, the shorter is its main-sequence lifetime
Fig II-59: Main Sequence Stars of Different Masses
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The Sun has been a main-sequence star for about 4.56 billion years
and should remain one for about another 7 billion years
Fig II-60: The Zero-Age Sun and Today’s Sun
During a star’s main-sequence lifetime, the star expands
somewhat and undergoes a modest increase in
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luminosity