# Lecture 14: The Sun and energy transport in stars. Astronomy 111

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1 Lecture 14: The Sun and energy transport in stars Astronomy 111

2 Energy transport in stars What is a star? What is a star composed of? Why does a star shine? What is the source of a star s energy?

3 Laws of Stellar Structure I: The Gas Law Most stars obey the Perfect Gas Law: Pressure = Density Temperature In words: Compressing a gas results in higher P & T Expanding a gas results in lower P & T

4 Laws of Stellar Structure II: The Law of Gravity Stars are very massive & bound together by their Self-Gravity. Gravitational binding increases as 1/R 2 In words: Compress a star, gravitational binding gets stronger. Expand a star, gravitational binding gets weaker.

5 Hydrostatic Equilibrium Gravity wants to make a star contract. Pressure wants to make a star expand. Counteract each other: Gravity confines the gas against Pressure. Pressure supports the star against Gravity. Exact Balance = Hydrostatic Equilibrium The star neither expands nor contracts.

6 Hydrostatic Equilibrium Gravity Gas Pressure

7 Core-Envelope Structure Outer layers press down on the inner layers. The deeper you go into a star, the greater the pressure. The Gas Law says: More pressure= hotter, denser gas Consequences: hot, dense, compact CORE cooler, lower density, extended ENVELOPE

8 Core-Envelope Structure Compact Core Extended Envelope

9 Example: The Sun Core: Radius = 0.25 R sun T = 15 Million K Density = 150 g/cc Envelope: Radius = R sun = 700,000 km T = 5800 K Density = 10 7 g/cc

10 The Essential Tension The Life of a star is a constant tug-of-war between Gravity & Pressure. Tip the internal balance either way, and it will change the star s outward appearance. Internal Changes have External Consequences

11 Why do stars shine? Stars shine because they are hot. emit thermal (~blackbody) radiation heat leaks out of their photospheres. Luminosity = rate of energy loss. To stay hot, stars must make up for the lost energy, otherwise they would go out.

12 Case Study: The Sun Question: How long can the Sun shine? Answer: Consider the internal heat content of the Sun. Luminosity = rate of heat loss. Lifetime = Internal Heat Luminosity

13 What if no source of energy? The Sun s Luminosity losses would not be balanced by input of new internal heat. The Sun would steadily cool off & fade out. 18th Century: Assumed a solid Sun (iron & rock) Found Lifetime ~ 10 Million Years No Problem: Earth was thought to be thousands of years old.

14 The Age Crisis: Part I Late 1800s: Geologists: Earth was millions of years old. How can Earth be older than the Sun? Astronomers: The Sun is a big ball of gas in Hydrostatic Equilibrium. Kelvin & Helmholtz proposed Gravitational Contraction as a source of energy.

15 Kelvin-Helmholtz Mechanism Luminosity radiates away heat Outer layers of the Sun cool at little, lowering the gas pressure. Lower Pressure means Gravity gets the upper hand and the star contracts a little. Contraction compresses the core, heating it up a little, adding heat to the Sun. Sun could shine for ~100 Million years

16 Gravity & Pressure in Equilibrium G P

17 Luminosity radiates away heat & Pressure Drops G P

18 Balance tips in favor of gravity, Sun shrinks. G P

19 Contraction makes core heat up, increasing the internal Pressure. G P

20 Balance restored, but with higher gravity, pressure & temperature than before... G P Starts the cycle all over again...

21 The Age Crisis: Part II Early 1900s: Geologists show that the Earth is >2 Billion years old. Kelvin Says: The Geologists are wrong. Nature Says: Kelvin is wrong... There is new physics.

22 Nuclear Fusion 1905: Einstein demonstrates that Mass and Energy are equivalent: E=mc s: Eddington noted that 4 protons have 0.7% more mass than 1 Helium nucleus (2p+2n). If 4 protons fuse into 1 Helium nucleus, the remaining 0.7% of mass is converted to energy.

23 Fusion Energy Fuse 1 gram of Hydrogen into grams of Helium. Leftover grams converted into Energy: E = mc 2 = 6.3x10 18 ergs Enough energy to lift 64,000 Tons of rock to a height of 1 km.

24 Hydrogen Fusion Question: How do you fuse 4 1 H (p) into 4 He (2p+2n)? Issues: Four protons colliding at once is unlikely. Must turn 2 of the protons into neutrons. Must be hot: >10 Million K to get protons close enough to fuse together.

25 Proton-Proton Chain 3-step Fusion Chain

26 neutrino 2 H positron photon 3 He 4 He 3 He 2 H positron photon neutrino

27 The Bottom Line Fuse 4 protons ( 1 H) into one 4 He nucleus. Release energy in the form of: 2 Gamma-ray photons 2 neutrinos that leave the Sun 2 positrons that hit nearby electrons, creating two more Gamma-ray photons Motions (heat) of final 4 He and 2 protons.

28 The Age Crisis: Averted Luminosity of the Sun is ~4x10 33 erg/sec Must fuse ~600 Million Tons of H into He every second. ~4 Million tons converted to energy per second. Sun contains ~10 21 Million Tons of Hydrogen Only ~10% is hot enough for fusion to occur. Fusion Lifetime is ~10 Billion Years.

29 Test: Solar Neutrinos Question: How do we know that fusion is occurring in the core of the Sun? Answer: Look for the neutrinos created in Step 1 of the Proton-Proton chain.

30 What are Neutrinos? Massless, weakly interacting neutral particles. Travel at the speed of light. Can pass through a block of lead 1 parsec thick! Any neutrinos created by nuclear fusion in the Sun s core would stream out of the Sun at the speed of light.

31 Solar Neutrinos: Observed! Detection of neutrinos is very hard: Need massive amounts of detector materials Work underground to shield out other radiation Answer: We detect neutrinos from the P-P chain in all the experiments performed to date! Success! But...

32 The Solar Neutrino Problem We detect only ~1/3 of the number expected... Why? Is our solar structure model wrong? Is the subatomic particle theory of neutrinos wrong? The last seems most plausible recent evidence neutrinos have mass question for the physics of 21st century!

33 Putting Stars Together Physics needed to describe how stars work: Law of Gravity Equation of State ( gas law ) Principle of Hydrostatic Equilibrium Source of Energy (e.g., Nuclear Fusion) Movement of Energy through star

34 Hydrostatic Equilibrium Balance between Pressure & Gravity. If Pressure dominates, the star expands. If Gravity dominates, the star contracts. Sets up a Core-Envelope Structure: Hot, dense, compact core. Cooler, low-density, extended envelope.

35 Hydrostatic Equilibrium Gravity Gas Pressure

36 Energy Generation Stars shine because they are hot. To stay hot stars must make up for the energy lost by shining. Two Energy sources available: Gravitational Contraction (Kelvin- Helmholtz) Only available if star is NOT in equilibrium Nuclear Fusion in the hot core.

37 Main-Sequence Stars Generate energy by fusion of 4 1 H into 1 4 He. Proton-Proton Chain: Relies on proton-proton reactions Efficient at low core Temperature (T C <18M K) CNO Cycle: Carbon acts as a catalyst Efficient at high core Temperature (T C >18M K)

38 Proton-Proton Chain

39 CNO Cycle:

40 Controlled Nuclear Fusion Fusion reactions are Temperature sensitive: Higher Core Temperature = More Fusion ε(pp) ~ T 4 ε(cno) ~ T 18 BUT, More fusion makes the core hotter, Hotter core leads to even more fusion, Why doesn t it blow up like an H-Bomb?

41 Hydrostatic Thermostat If fusion reactions run too fast: core heats up, leading to higher pressure pressure increase makes the core expand. expansion cools core, slowing fusion. If fusion reactions run too slow: core cools down, leading to lower pressure pressure drop makes the core contract. contraction heats core, increasing fusion.

42 Thermal Equilibrium Heat always flows from hotter regions into cooler regions. In a star, heat must flow: from the hot core, out through the cooler envelope, to the surface where it is radiated away as light.

43 Energy Transport There are 3 ways to transport energy: 1) Radiation: Energy is carried by photons 2) Convection: Energy carried by bulk motions of gas 3) Conduction: Energy carried by particle motions

44 Radiation Energy is carried by photons. Photons leave the core Hit an atom or electron within ~1cm and get scattered. Slowly stagger to the surface ( random walk ) Break into many lowenergy photons. Takes ~1 Million years to reach the surface.

45 Convection Energy carried by bulk motions of the gas. Analogy is water boiling: Hot blob rises cooler water sinks

46 Conduction Heat is passed from atom-to-atom in a dense material from hot to cool regions. Analogy: Holding a spoon in a candle flame, the handle eventually gets hot.

47 The Sun

48 Energy Transport in Stars Normal Stars: A mix of Radiation & Convection transports energy from the core to the surface. Conduction is inefficient (density is too low). White Dwarfs: Ultra-dense stars Conduction dominates energy transport.

49 Summary: Stars shine because they are hot. need an energy source to stay hot. Kelvin-Helmholtz Mechanism Energy from slow Gravitational Contraction Cannot work to power the present-day Sun Nuclear Fusion Energy Energy from Fusion of 4 1 H into 1 4 He Dominant process in the present-day Sun

50 Summary: Energy generation in stars: Nuclear Fusion in the core. Controlled by a Hydrostatic thermostat. Energy is transported to the surface by: Radiation & Convection in normal stars Conduction in white dwarf stars With Hydrostatic Equilibrium, these determine the detailed structure of a star.

51 Summary: Stars are held together by their selfgravity Hydrostatic Equilibrium Balance between Gravity & Pressure Core-Envelope Structure of Stars Hot, dense, compact core cooler, low-density, extended envelope

52 Lecture 8: Telescopes Refracting vs. reflecting Angular resolution of a telescope Properties of different types of telescopes What is seeing? What are CCDs?

53 Lecture 9: Formation of the Solar System How did solar system form? Characteristics of inner and outer Solar System How did angular momentum play a role? Define: planetesimals, radioactive age dating, terrestrial, jovian Name the four techniques astronomers use to find extrasolar planets

54 Lecture 10: Terrestrial planets Characteristics of Earth, the Moon, and terrestrial planets How do we know the age of the Earth? Where did the Moon come from?

55 Lecture 11: Jovian planets Characteristics of Jovian planets, dwarf planets, and other solar system objects Define dwarf planet, asteroid, Kuiper belt, comet

56 Lecture 12: Distances to stars Calculate distance to a star given its parallax Define: apparent brightness, luminosity, true binary, visual binary, spectroscopic binary, eclipsing binary What quantities can we measure using binary stars?

57 Lecture 13: Colors of stars Why are stars different colors? Why do stars have absorption lines in their spectra? Give the order of the spectral sequence Be able to calculate luminosity and know its units (energy/sec) Be able to use the Stefan-Boltzman law and know its units (energy/sec/area) Know the main components and axes of an H-R diagram

58 Lecture 14: Energy transport in the Sun Define: hydrostatic equilibrium, P-P chain, CNO cycle, neutrinos Know what the Perfect Gas Law means Describe how the P-P chain makes energy in the Sun through fusion List the three modes of energy transport

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