Download Lecture 16, PPT version

Outline - March 25, 2010
• Recap: Evolution and death of low mass stars (pgs. 566-572)
• Evolution and death of high mass stars (pgs. 572-581)
• Stellar Remnants (white dwarf, neutron star, pulsar, black hole)
• Novae and supernovae
H-R Diagram
About 90% of stars in the sky
are “Main Sequence” stars
All main sequence stars are
stable (gravity exactly
balances pressure) and
energy source is fusion of
HYDROGEN to form
HELIUM
All of the “non-main
sequence” objects are no
longer burning H in their
cores (are “evolved” stars)
Build-up of Inert Helium Core
Eventually, the star builds up a substantial He
core, with H “burning” in a shell around the
core.
The H burns into layers of the star that are
thinner, and thinner, making it harder to
hold the star up against gravitational
collapse.
The He core can provide a little bit of help
by contracting (conversion of gravitational
energy, just like a protostar).
As the core contracts, the outer envelope
expands and the star leaves the main
sequence.
Evolution of a Low-Mass Star
As He core contracts, the star moves up the
HR diagram.
As outer envelope expands, the star becomes
physically larger (increases luminosity) and the
surface temperature cools (becomes redder).
Star becomes a Red Giant.
Onset of He burning in the core happens quite
suddenly (helium “flash”) once the
temperature and density of the core are
high enough to fuse He.
Helium flash doesn’t disrupt the star
(localized region of 1/1000 of the star), but
does cause the core to expand a little bit
(and envelope shrinks in response).
Red Giant Phase for Low Mass Stars
Core is now 100 million Kelvin (about 10x
hotter than when the star was a main
sequence star)
Two sources of energy:
1. H to He in a shell
2. He to C in the core
Final Stage of Evolution of Low-Mass Star
It’s only a matter of time before the star gets in trouble again…
This time it’s CARBON ash that has sunk to the center (non-burning
carbon core, surrounded by a shell of He burning, surrounded by a
shell of H burning).
Most low mass stars can repeat the core contraction process, and
ignite Carbon fusion (which produces Oxygen).
But, once a significant amount of oxygen has built up in the
core, it’s game over for the star!!
Death of a Low Mass Star
Carbon-Oxygen core contracts in an attempt to help hold the star up against
gravitational collapse; but there isn’t enough mass in the star to make the
temperature and density high enough to fuse the oxgygen
Core shrinks down to about the size of the earth, and can’t go any farther
because of a quantum mechanical effect
Can only compress electrons so far - this is what stops the core contraction
Pressure in the core is provided by “degenerate electron gas” and the core
becomes stable (no longer contracting)
“Burning fronts” (H, He, C) plow out into the very light, fluffy layers of the
(enormous!) star, and the outer layers of the star “lift off” due to radiation
pressure
Formation of White Dwarf and Planetary Nebula
(end of a low-mass star)
Outer layers of star lift off, revealing small, hot core = White Dwarf
and creating a “Planetary Nebula”
Sirius B (white dwarf
companion to Sirius A)
Evolutionary Track on the HR Diagram
(Low-Mass Star)
Evolution of High-Mass Stars
Unlike low mass stars, high mass
stars make a steady transition
from H fusion in the core to He
fusion in the core (no “helium
flash”), to O fusion in the core,
and they keep on going to heavier
chemical elements.
High-mass stars evolve off the
main sequence to become
“supergiant” stars.
“Onion Layers” of Fusion in a High-Mass Star
Star undergoes cycles of
core contraction and
envelope expansion,
fusing heavier and
heavier chemical
elements, until an iron
core forms.
Once silicon starts to
fuse, the star has about
a week to live.
Timescales of Fusion
(Mstar = 20 Msun)
H fusion in core: 10 million years
He fusion in core: 1 million years
C fusion in core: 1000 years
O fusion in core: 1 year
Si fusion in core: 1 week
What’s so special about Iron (Fe)?
Fusion of nucleii that are lighter than iron result in a net gain of
energy (takes less energy to bring the nucleii close together than
you get from mass loss)
Fusion of nucleii that are as heavy or heavier than iron result in
a net loss of energy (takes more energy to bring the nucleii
close together than you get from mass loss)
Bottom line: star can’t use iron as a nuclear fuel to support
itself from gravitational collapse, because fusing iron is a losing
proposition in the energy balance!
Death of a High-Mass Star
Supernova: Implosion followed by Explosion
•
Once substantial amount of iron has built up, star implodes on itself
•
Core reaches temperature of 10 billion Kelvin (= tremendously high energy
photons), the nuclei are split apart into protons and neutrons
(“photodisintegration”)
•
In less than 1 second, the star undoes most of the effects of nuclear fusion that
happened in the previous 11 million years!!!!!
•
High-energy photons are absorbed, giving rise to loss of thermal energy in the core,
the core becomes even more unstable, and the collapse accelerates
•
Protons and electrons in the core combine together (“neutronization”), resulting in
nothing but neutrons in the core
•
Collapse continues until it’s not possible to squeeze the neutrons together any
tighter (size of core = size of Manhattan)
•
Collapse starts to slow, but overshoots and outer layers of star are driven out
into space (perhaps by “bounce” off the neutron core) in a massive explosion
Supernovae Generate Tremendous Amounts of Energy
At their maximum
brightness, supernovae
are as bright as an
entire galaxy.
Peak luminosity is
about 1051 ergs = the
sun’s total output of
energy over 10 billion
years!
How long does a supernova last?
“Type II” supernovae
are exploding highmass stars
“Type Ia” supernovae
are something else
entirely (and involve
binary star systems)
Why should you care about supernovae?
• Extraordinarily bright, so can use them to measure distances
to galaxies that are very far away: b = L / (4 d2)
• Supernovae are the source of all heavy chemical elements!
• The heavy chemical elements are produced during the
explosion itself, when there is more than enough energy to
fuse nuclei heavier than iron (doesn’t matter that there is a
“net loss” of energy - the star is already VERY far out of
equilibrium)
Supernova Remnants
(high-mass star guts)
Cycle of Star Formation and Supernovae
• Stars form out of gas in the ISM, evolve, and blow much of
themselves back into the ISM
• Massive stars create heavy chemical elements during the
explosions, which “enriches” the ISM with heavy chemical
elements
• New stars form, and make yet more heavy chemical elements
• It takes about 500 cycles of massive star formation to account
for all the heavy chemical elements in the universe
• More than enough time for this to happen (universe is 14
billion years old, massive stars take a few million years to evolve
and explode)
Stellar “Remnants”
What’s left behind after a star dies?
Main sequence mass < 5 Msun: white dwarf
Main sequence mass between 5 Msun and 40 Msun: neutron star
Main sequence mass > 40 Msun: black hole
All of these are stable (neither expanding nor contracting), so
long as they are “left alone”.
Pressure in white dwarf and neutron star is somewhat exotic
(not normal gas pressure or radiation pressure) due to their
highly-compressed states.
White Dwarfs
• Pressure comes from “degenerate electron pressure”
• Electrons packed together as tightly as quantum mechanics allows;
their speeds support the WD against gravitational collapse
• WD acts a lot like a metal (same temperature and density throughout)
• Maximum WD mass = 1.4 Msun (“Chandrasekhar limit”)
• WD with mass = 1 Msun is about the size of the earth, weight of 1
teaspoon of WD material = about the weight of a small truck
• If all alone in space, WD simply cool off (no internal source of energy)
and eventually become black
White Dwarfs in Binary Systems
•
Most stars are found in binary
systems
•
May have situation where WD orbits
a giant or supergiant star at a
relatively close distance
•
Outer layers of the giant or
supergiant are very light and fluffy,
and may be pulled over onto the
WD by gravity
•
Material from companion star builds
up in an “accretion disk” around the
WD, and eventually winds up on the
surface of the WD
White Dwarfs in Binary Systems, II
What happens to the WD when mass is dumped onto it
depends on how much mass, and how fast.
Slow accretion of not much mass (not enough to make the
mass of the WD > 1.4 Msun): nova
Fast accretion of a lot of mass (enough to make the mass of
the WD > 1.4 Msun rather suddenly): supernova (“Type 1a”)
Novae
• Thin layer of (mostly) H from the companion star builds up on
surface of WD
• Sudden flare in brightness (increases by about a factor of
10,000 or more), then fades over the course of about a month
• Flare is due to hydrogen fusion on the surface of the white
dwarf
• Novae happen about 2 or 3 times per year in our Galaxy
• Can recur (i.e., same WD can “go nova”, but not very
predicable)
Novae
H fusion on the surface of a WD
“Naked eye” nova; picture taken
in the Varzaneh Desert in
Isfahan, Iran (February 2007)
White Dwarf Supernova
Supernova Type Ia
If the companion star to a WD dumps a lot of mass onto the WD very quickly, making the
mass of the WD exceed the Chandrasekhar mass (1.4 Msun), the WD explodes as a
supernova!
WD is much like a hot metal ball, same temperature and same density throughout
Addition of extra mass causes WD to contract (gravity “wins” over pressure from the
electrons) and instantaneously the carbon starts to fuse throughout all parts of the WD,
blowing the WD to bits
Two Basic Types of Supernovae
Note: Supernovae NEVER repeat!
Remnants of two different supernovae.
Left: a Type Ia supernova (WD).
Right: a Type II supernova (high mass
star). This is a “happy alignment” of
images - the two stars weren’t related
to each other!
What’s left behind after a massive star goes supernova?
If the mass of the core is less than about 3 Msun, a neutron
star is left behind.
If the mass of the core is greater than about 3 Msun, there is
no source of sufficient pressure to keep the core from
collapsing completely under gravity, and a black hole is
formed.
Neutron Stars
•
Even more compressed than WD
•
Typically the size of a city (about 10 km in radius) with mass between 1.4 Msun
and 3 Msun
•
Density is such that the weight of one teaspoon of NS material would weigh 100
million tons (vs. 1 ton for WD material)
•
NS supported by “neutron degenerate pressure” (again, quantum mechanical
phenomenon having to do with how tightly neutrons can be packed together)
•
Must rotate extremely fast (conservation of angular momentum); a star that was
originally rotating once per month would now have to rotate a few times per
second!
•
Compression of the material also compresses the magnetic field, and amplifies
its strength (making it trillions of times larger than the earth’s magnetic field)
Pulsars
Rapidly-Rotating Neutron Stars
First discovered by Jocelyn Bell (1967) as “pulses”
of radio light coming from the Crab Nebula
The pulses lasted 0.01 seconds, and repeated every
1.4 seconds
In 1974, Jocelyn’s PhD advisor (Tony Hewish) got the
Nobel Prize for explaining what puslars are
Now know of 100’s of pulsars, most with periods
between 0.03 and 0.3 seconds (meaning they rotate
between 3 and 30 times per seconds)
The fastest known pulsars have periods of
milliseconds and are rotating at speeds approaching
0.25c !!!!
The radio light from plusars is called “synchrotron radiation”, a type of light
that is emitted by electrons as they move on spiral paths around magnetic field
lines. The “synchrotron radiation” pulses are proof of the fast rotation rates of
neutron stars and the presence of an incredibly strong magnetic field.
Pulsar at Center of Crab Nebula
Supernova observed by Chinese astronomers in year 1054
Crab Nebula - remnant
of supernova explosion
Pulsar Recordings
http://www.astrosurf.com/luxorion/audiofiles-pulsar.htm
Crab - 1.4 rotations per second
Vela - 11 rotations per second
PSR 1937+21 (a
“millisecond” pulsar) 642 rotations per second
Midterm Exam #2
Curve boundaries for Midterm #2:
A
> 90.5%
A-
87.5% to 90.5%
B+
81% to 87.5%
B
75% to 81%
B-
69% to 75%
C+
66% to 69%
C
63% to 66%
C-
57% to 63%
D
50% to 57%
F
< 50%
Class letter grade average
based on the curve is
between B and B- (2.9 / 4.0)
Your score on this exam: 78.5 / 100
Your ranking in the class on this exam:
21 / 47
Approximate letter grade on this exam: B
This info is on the last page
of your exam.
Approximate Mid-semester Grades
A
> 92%
A-
86% to 92%
B+
82% to 86%
B
78% to 82%
B-
70% to 78%
C+
66% to 70%
C
62% to 66%
C-
57% to 62%
D
50% to 57%
F
< 50%
Approximate mid-semester grades based on average of midterm
exams, best 3 of 4 home work assignments, and average of 2 labs.