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Star Formation Daniel Zajfman Department of Particle Physics Weizmann Institute of Science.

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Presentation on theme: "Star Formation Daniel Zajfman Department of Particle Physics Weizmann Institute of Science."— Presentation transcript:

1 Star Formation Daniel Zajfman Department of Particle Physics Weizmann Institute of Science

2 Stars Planets Galaxies Black holes Nebulae Red giants White dwarf Moons Supernovae Pulsars Neutron stars Why so many different objects? Why all the stars are not alike?

3 Stars are not permanent objects: They are born, live and die, just like human being

4 Big Bang Nucleosynthesis Mainly Hydrogen, Deuterium and Helium  Star should “work” with these materials

5 Elements and Isotopes We define an “element” by the number of protons in its nucleus. There can be “isotopes” with different numbers of neutrons.

6 Big Bang Particle Physics Nuclear Physics Matter-Radiation Equilibrium Atomic & Molecular Physics H + +e -  H+hν “Recombination era” Pre-galactic gas clouds First generation of stars 10 12 K 5x10 9 K 5x10 8 K 4x10 3 K 4 He, D, 3 He, 7 Li 100 s 1000 s 10 6 years Time scale Temperature

7 The Universe after the Big-Bang is “uniformly” filled with Hydrogen, Deuterium, and Helium Small fluctuations (finite number of particles) create small lump of matter, which start to collapse under their own gravity Formation of protogalaxies

8 Anatomy of an interstellar cloud Collapse of molecular clouds: Not in a single piece (clumps formation) Clumps collapse to form stars 10-1000 stars can be formed from one single cloud Mostly in second generation clouds

9 Horsehead Nebula Barnard68 Eagle Interstellar clouds are the nursery of stars. Some clouds, called molecular clouds, contain a minor (but important) fraction of molecular species. RCW38

10 The beginning: The birth of a star

11 Cloud collapse Method 1 Build up of small clouds to giant ones Clouds stick together and grow Gravitation takes over Very slow process (low interstellar density) Method 2Gravity and radiation pressure Method 3 Compression by supernova blast waves Not for first stars!

12 Two hindrances to collapse Internal heating: Potential energy  Kinetic energy (Gas particles speed up and collide) Temperature increases Pressure build up which slows (or stops) the collapse Energy is radiated away Angular momentum L=mass x vel. of rotation x radius (L=mvr) Conservation of angular momentum: Constant for a closed system Thus, as the cloud shrinks due to gravity it spins fasters  Collapse occurs preferentially along path of least rotation  The cloud collapses into a central core surrounded by a disk Gravity makes the cloud collapse!

13 Orion cloud ~1000 ly Proplyds Protostar and Proplyds

14 Planet formation??? Protostar

15 The process can be very “unstable” and often yields to the production of “jets” for about 100,000 years

16 Protostars and jets

17 Protostar formation The central core is called a protostar Surface ~ 300 K, the internal temperature is steadily increasing Undergoing continuous gravitational contraction Self-compression heats the central core Nuclear Fusion reaction starts A star is born

18 Planets are probably formed later in the remaining disk of the protostar

19 A more detailed look at the collapse process allows to extract the critical mass of a cloud so that a star can be formed Sir James Jeans: the critical mass, called today the Jeans mass Can we estimate it easily?? R Let’s assume we compress the gas slightly. It will bounce back to its original size in a time At the same time, the gravity will attempt to contract The system, and will do that in a “fall-free” time G is the universal gravitational constant ρ is the gas density If we want Gravity to win, we need: Jeans Mass

20 Star formation – The movie

21 Gas cloud Young stars 100 million years The size of the cloud changes from million of km down to few thousands of km. The temperature increases from -270 o C to million of degrees. At this temperature, the nuclear fuel (hydrogen) is “light up”. Light is emitted, and the star starts its life: on one hand, gravitational force pushes inward, while on the other hand internal pressure due to the nuclear reactions pushes outward.

22 Thermonuclear Fusion In order to get fusion, one must overcome the electric repulsion. You can do this by having high density (lots of particles) and high temperature (particles moving very quickly).

23 For Stars, size matters A star mass determines which fusion reaction are possible in the core, and hence its luminosity, surface temperature and lifetime. Object with mass smaller than 8% of the solar mass (75 times Jupiter mass) never ignite fusion, and therefore fade to obscurity in about 100 million years. These are Brown Dwarf. Sun mass: 2 x 10 30 kg Jupiter Mass: 2x10 27 kg First ever observed brown dwarf in October 1994 How many brown dwarf in the Universe?

24 The sun: A typical star Age: ~ 4-5 billion years old

25 The Power of the Star: The Proton-Proton Cycle This is the primary source of energy for main sequence stars Minimum temperature: 5 millions K In this reaction cycle, 4 protons are transformed in one He nuclei, 2 positrons, gamma rays and 2 neutrinos

26 Another view of the proton-proton cycle Each reaction cycle requires 4 hydrogen (protons) and yields about 25 MeV of energy

27 The proton-proton cycle is the most important reaction in the sun. Is there enough Hydrogen? Let’s estimate the lifetime of the sun In the p-p cycle, each time 4 protons react, and produce one 4 He nuclei 4p  4 He + 2e + m p =1.67x10 -27 kg m He =6.6326x10 -27 kg m e+ =9.1139x10 -31 kg 4m p =6.68x10 -27 kg m He +2m e+ =6.6344x10 -27 kg Mass difference: ∆m=4.56x10 -29 kg Where did this mass goes?? E=∆mc 2 !! How much energy is thus produced in one p-p cycle? E=∆mc 2 = 4.56x10 -29 kg x (3x10 8 ) 2 (m/s) 2 = 4.1 x 10 -12 Joule That’s by the way 25 MeV!

28 We know that the total power output of the sun is: L=3.9 x10 26 Joule/second (eq~ 100 billion nuclear bomb/second). Lifetime of the sun (cont.) Thus, the number of p-p cycle per second in the sun is: Total power/energy per cycle=L/E=3.9x10 26 /4.1x10 -12 =9.5x10 37 reactions/second Since each p-p cycle requires 4 protons, the number of protons used every second in the sun is: n p =4x9.5x10 37 =3.8x10 38 protons/second How many protons are in the sun? #protons~ mass of sun/mass of protons = 2x10 30 kg/1.67x10 -27 kg ~ 1x10 57 protons Thus, the lifetime of the sun is approximately: 1x10 57 /3.8x10 38 =2x10 18 seconds which are about 60 milliard years. However, the sun uses only 10% of its hydrogen… so lifetime is of the order of (very roughly) 6 milliard years

29 The CNO Fusion Cycle For more massive stars (higher temperature) In this cycle, 4 protons are converted into 1 Helium, 2 positrons, gamma rays and 2 neutrinos Why more massive stars? Because of the electrostatic repulsion of the Carbon nuclei In the sun, this produce only 2% of the total energy!

30 The triple alpha process Three Helium nuclei are converted into a carbon nucleus and gamma rays For star leaving the main sequence (called Red Giants) Nucleosynthesis!

31 Comparison of the p-p and CNO cycle Usually the CNO cycle is more important for heavier stars, as it is hotter inside

32 The lifetime of a star depends (mainly) on its mass Higher the mass, shorter the lifetime! High mass: M > 8M sun Intermediate mass 2M sun < M < 8M sun Low mass: M { "@context": "", "@type": "ImageObject", "contentUrl": "", "name": "The lifetime of a star depends (mainly) on its mass Higher the mass, shorter the lifetime.", "description": "High mass: M > 8M sun Intermediate mass 2M sun < M < 8M sun Low mass: M

33 Can we prove (experimentally) that all that is correct? The Solar Neutrino (ex)-Problem If the sun is really powered by nuclear (fusion, p-p cycle) power, then it has to produce some special particles called neutrinos. “The” solar neutrino problem: "the sun does not produce enough neutrinos" These particles have almost no interactions with matter, get out of the sun core, and can be detected by terrestrial neutrino detectors.

34 The 37 Cl neutrino detector is a tank containing 375,000 liters of Perchloroethylene in a cavity 1,500 m below ground When a neutrino (with the right energy) collides with a 37 Cl atom, it produces an atom of 37 Ar (and an electron) which is radioactive, and can be detected later. Ray Davis, 1966 Nobel prize in 2002 First experiment in Homestake mine Only ~ 1/3 of the expected Neutrino were measured

35 The super Kamiokande detector Detecting neutrino coming from the center of the sun. Produce the first evidence (1998) that something was “wrong” with the neutrino physics

36 The final word It is the physics of neutrino which was “wrong” Neutrino oscillation The solar fusion theory is correct Neutrino have masses (Particle Astrophysics is a very rich and exciting field of Physics)

37 The Hertzsprung-Russel (HR) Diagram Star Characterization Reversed scale!!!

38 When one plot the data of a group of star (for example close to us) This is what we see on the HR diagram What is the “Main Sequence”?? The HR diagram

39 Temperature, Size and Luminosity Hotter objects are brighter Energy radiated per unit of time and unit of area is proportional to T 4 Thus, larger Temperature means more energy radiated Bigger objects are brighter Energy radiated per unit of time and unit of area is proportional to T 4 Thus larger surface means more energy radiated Let’s assume all stars are the size of the Sun, but the hotter ones are more luminous, just because they are hotter Then all the stars would fall on the blue line In math-language it means: Surface Stefan-Boltzman Law

40 In reality: Not really true! But we learned something: The coolest main sequence stars are a lot smaller than the sun. The hottest main sequence stars are a lot bigger than the sun.

41 The Hertzsprung-Russel (HR) Diagram Spectral classes instead of temperature Our sun is spectral class G

42 In general, the HR diagram allows to categorize the different stars using “measureable” parameters. Different type of stars are located in different region of this diagram.

43 M5 cluster with more data points and a calculated isochrone line The line represent the calculated “behavior” of a star in the H-R diagram assuming all stars have the same age (but were born with different initial size) The Best Physics we know today is in good agreement with observations

44 Stellar Lifetimes

45 Next episode: Stellar evolution Nucleosynthesis Binary systems Final stages Supernovae Black Holes Quasars Pulsars Interstellar medium

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