Astronomy 106 Review for Exam 1

Transcription

1 Astronomy 106 Review for Exam 1 A Review Of: Age Dating Techniques The Drake Equation Chemical Composition of Life Structure of the Atom Nuclear Fusion Generations of Stars Star/Planet Formation Temperature of Planets Habitable Zone Detecting Exoplanets Tidal Forces Ryan Rubenzahl Undergraduate Teaching Assistant Department of Physics and Astronomy September 30th 2015

2 How do we know the age of things? Age of the Universe? Calculate from the expansion caused by the Big Bang Observe galaxies moving apart, measure their velocities Run time backwards until all galaxies are in same spot 13,700,000,000 years old Age of the Milky Way? Can t be older than Universe, so at most 13.7Gyr Look at oldest objects in Milky Way to find minimum age White-dwarfs form from stars that live for about 4Gyr Oldest White-dwarfs observable are ~9.5Gyr Hubble Ultra-Deep Field Star 4Gyr White-dwarf 9.5Gyr Now Total time = 4Gyr + 9.5Gyr = 13.5Gyr, So the Milky Way is ~13.5Gyr

3 How do we know the age of things (cont.)? Living things contain radioactive Carbon-14, which doesn t decay while the organism is alive. Once it dies, the C-14 decays into the rare (and stable) Nitrogen-14: C yrs C-14 N yrs C-14 N-14 The time it takes for half the initial amount to decay is a known rate, called the half life. For Carbon, this is 5700 years. This value is unique to the element that is decaying. amount remaining 1 n = n o 2 original amount t t 1/2 time thing has been decaying for half life Rocks aren t alive, but they do contain different radioactive elements

4 How do we know the age of things (cont.)? Age of the Solar System? C-14 doesn t work well for dating rocks For one, C-14 is only found in living or once-living things, and rocks were never alive C-14 s half life is too short to date rocks effectively luckily rocks contain a different element: Rb-87 49,900,000,000yrs Does not participate in radioactive decay process half-life Rb-87 Sr-87 Sr-86 Radionuclide Daughter Reference N = D = n(radionuclide) n(reference) n(daughter) n(reference) = n(rb-87) n(sr-86) = n(sr-87) n(sr-86) Given two minerals A, B t1 2! D ln A DB " t = $ + 1 % ln 2 & NA NB '

5 How do we know the age of things (cont.)? The halflife of 87 Rb is measured to be 10 t 1 2 = years. Example. Samples of two minerals from the same igneous rock from northern Ontario are analyzed in a mass spectrometer, with these results for the number ratios N and D: Mineral Rb Sr(N) Sr Plagioclase (B) Pyroxene (A) Sr (D) How old is the rock? t t1 2! D ln A DB " = $ + 1 % ln 2 & NA NB '

6 Age Summary So we have these experimental facts: The Universe is 13.7 billion years old, give or take about 0.1 billion. The Milky Way Galaxy is about 13.5 billion years old; certainly it cannot be younger than the oldest white dwarfs it contains, which are 10 billion years old. The Solar system Sun, planets, asteroids, etc. is billion years old, give or take about a million years. The Earth s surface solidified about 3.8 billion years ago. These timespans are much longer than the age the world was thought to have in Darwin s time. This has expanded dramatically the scope of the slow processes of evolution.

7 Drake Equation

8 Chemical Composition of Life As We Know It Much of LAWKI is similarly composed, but differs from the Earth s composition and is in fact more closely similar to the Solar System at large. By atom, per hundred oxygen atoms, Hydrogen Carbon Nitrogen Oxygen Element Bacteria Human Interior Crust Ocean Air Earth Sun Comets Phosphorus < < Sulfur Calcium <

9 Structure of the Atom Atoms are composed of three particles, or ingredients Proton, + electric charge Neutron, 0 electric charge Electron, - electric charge N P A Hydrogen Atom E Electrons are held in orbit around nucleus by the Electrostatic Force Protons + Neutrons are held together by the Strong Nuclear Force - Strong force only acts on very short distances (radius of proton) - In order to stick protons together, they must be pushed extremely close - This requires extremely high temperatures and pressures

10 Atomic Energy Levels Atoms have fixed energies corresponding to specific Energy Levels - Electrons in the atoms can move between Energy Levels under right circumstances - If an electron gains energy, it can move up a level. If it loses energy, it will move down a level E3 E2 E E3 E2 E Outgoing light with E = E2 - E1 E1 E Incoming light with E = E2 - E1 E1 E Every atom has a unique set of energy levels, so by measuring the spectra, or set of energy lines, of an object emitting light, we can determine its composition. - The relative brightness of lines can be used to determine density, temperature, and pressure of the emitting region.

11 Nuclear Fusion If we squish two atoms together at very high temperature and pressure we can create a new atom. This process can release A LOT of energy. A simplified example: Hydrogen + Hydrogen + Hydrogen + Hydrogen = Helium + Energy But the amount of energy produced decreases as the mass of the elements being fused increases. After Iron-56, the fusion reaction no longer produces energy but instead requires energy. There are two known places in which the conditions for fusion exist: 1. The centers of stars (stellar-core nucleosynthesis) 2. Blast waves from exploding stars or the Big Bang itself (explosive nucleosynthesis) The Universe started with all its mass in light elements and worked its way up through these two methods.

12 Big Bang Nucleosynthesis Within 1 second the Universe had cooled to the point where protons, neutrons, and electrons could form without being destroyed instantly. These particles could combine to become the first Hydrogen atoms After about 100 seconds these Hydrogens started fusing into Helium atoms Other than very small amounts of Lithium and Beryllium-7, no other elements were made during the Big Bang For heavier elements, we need prolonged, sustained fusion reactions Thankfully, stars provide just that

13 The First Stars Called Population III Stars, these first stars formed around 300 million years after the Big Bang, creating the Universe s first galaxies (including ours, remember the age we found for the Milky Way before?) These first stars were HUGE, typically times as massive as the Sun Modern stars are typically times the mass of the Sun The Crab Nebula, a supernova remnant (Hubble Space Telescope) Despite having so much mass, these stars burned through their fuel extremely rapidly (within a few million years) These stars were able to fuse elements up to Nitrogen (A=7) before dying When these stars died, they collapsed on themselves and exploded in what is called a corecollapse supernova These explosions created conditions capable of producing elements beyond Nitrogen and spread these ingredients into interstellar space

14 The Second Stars These stars formed from the leftover guts of the dead Population III stars The guts now contained heavier elements, including the Universe s first Carbon Much smaller than previous stars = MUCH longer lived, dimmer stars New elements produced from fusion in these stars cores, and additional nucleosynthetic options New ways to die, as seen from planetary nebulae and new types of supernova Most of these stars are no longer around, but some of the smaller ones still exist Globular cluster M13, in true color (UR/Mees Observatory). Contains about 500,000 stars. These Population II stars mostly now live in globular clusters, all ~13Gyr old Those larger than 10 solar masses still died in supernova Smaller ones lived at most a few billion years, but died in a much less spectacular way

15 Stars can now begin to burn Helium, which is a mechanically unstable process Over time, the outer layers of the star begin to swell up (making the star a giant or supergiant) Inner material starts to mix with the outer layers of the star The star sheds its outer layers, releasing the heavy-element enriched material into interstellar space After a few million years all that is left is the dead core of the star, a White Dwarf A New Way to Die White Dwarfs can actually re-die if they are part of a multi-star system One star in system dies, becomes white dwarf, and begins accreting, or taking, matter from companion star If the amount of matter accreted pushes the white dwarf beyond 1.4 solar masses, rapid collapse occurs and the star explodes in a Type IA Supernova The Helix Nebula, a planetary nebula, the outer layers shed by a deceased star (NASA, ESA, Vanderbilt University)

16 Population I Higher abundances of heavier elements Orbits confined more tightly to plane of Milky Way, not random like Pop II Absent from globular clusters, but can form in open clusters The Sun is one of these stars Many successive generations of stars have enriched the interstellar medium with wide ranges of light/heavy elements: First via explosive nucleosynthesis Later with stars s internal fusion products The interstellar medium is composed of 99% gas (by mass) and 1% dust This mixture leads to the formation of lots of molecules Our sun, a population I star (NASA)

17 Molecules in the Interstellar Medium How do molecules form from atoms in the ISM? Three ways: Dust grain catalysis: H atom colliding and sticking with dust grain, moving along surface, binding with another stuck H atom, recoiling/kicking dust grain away Ion-molecule reactions (zero threshold) Neutral-neutral reactions in shocked material (high threshold) Why isn t all the ISM in molecular form? Ultraviolet starlight destroys molecules, when they re unprotected by lots of gas and dust The ISM of the Milky Way has a mass of: 90% diffuse clouds 10% molecular cloud complexes These molecular clouds can collapse gravitationally to form new stars/solar systems Eagle Nebula, starforming region (NOAO)

18 Star Formation As molecular clouds collapse, molecules within the cloud collide and heat up These fragments are slowly tumbling to start with, and collisions cause them to tumble more As the cloud collapses, this rotational speed increases due to conservation of angular momentum (think of figure skater pulling their arms inward) The material in the center of the cloud, where the collapse is heading, becomes very dense and hot The rotating cloud flattens around the dense core into an accretion disk, as particles collide and fall inwards, heating the core even more Once the core becomes hot enough to sustain fusion, it ignites into a newborn protostar Over time, the disk dissipates due to processes that use up or drive away the surrounding gas After ~3-5 million years almost all tiny dust particles have been removed from the disk

19 Protoplanetary Disks Not all of the debris makes it into the central star Particles in the accretion disk can clump and form larger objects Bigger clumps are more likely to bump into other particles, and also have stronger gravity which sucks up nearby particles The original gas cloud is not necessarily uniform in density - gravitational instabilities may exist Gravitational instabilities in the gas in the disk help to create gas giant planets Gravitational instabilities in the dust in the disk help to create rocky terrestrial planets As the planets grow in size, their gravity sweeps up everything nearby in a runaway core-accretion This cleans up the planet s orbit of most other debris nicely, leaving just the planet behind HL Tau protoplanetary disk (Atacama Large Millimeter Array)

20 Temperature of Dust Grains Planets can be thought of as very large dust grains All opaque objects emit light The Flux (power per given area) emitted by a body with temperature T is given by F = σt where σ = erg sec cm K. Stefan s Law Stefan-Boltzmann constant So how to find the temperature of dust grain orbiting a star? Using Stefan s Law and the geometry of the orbit, we find that T = 279K Fast Rotator: 0 1 (1 A)L[L J A (r[au]) 2 1/4 Slow Rotator: 0 1 (1 A)L[L J ] T = A (r[au]) 2 1/4 Where A is the albedo of the dust grain, L is the luminosity of the star in solar luminosities, and r is the orbital radius of the grain in Astronomical Units

21 The Habitable Zone Star Too Hot! Just Right Too Cold! 0 expression for the orbital radius: T=373K 1 A T=273K If water boils at 373K and freezes at 273K, where does a planet need to be to be in the range where water is in its liquid state? Rearranging the temperature equation from the previous slides gives us an Fast Rotator: Slow Rotator: J 1 A r[au] = 279K T 2 q (1 A)L[L J ] r[au] = 394K T 2 q (1 A)L[L J ] We can use this equation with the temperature at the inner and outer boundaries of the habitable zone to see where planets could have liquid water on their surface

22 Observing Exoplanets Exoplanets are difficult to find - imagine looking at our Solar System from far away. The Sun would appear 2 billion times brighter than Jupiter! So how do we find them? Three main methods for detecting exoplanets: 1. Direct Imaging: Take pictures over a period of time and actually watch the planet orbit its host star 2. Radial Velocity (RV): Measure tiny, periodic wobble in star s motion along motion of Doppler shift 3. Transits: Periodic eclipsing of star by planet, or vice versa Of today s exoplanets, 62% have been detected by transits, 31% by RV, and 3% by imaging. The remaining 4% were identified by a variety of lesser-used means.

23 Observing Exoplanets (cont) Combining the transit method with RV can provide even more information about an exoplanet: Aquire radius from transit, mass from RV, from radius we can calculate the volume of the planet and combined with the mass we can get a density This allows us to determine whether a planet is gaseous like Jupiter or rocky like Earth Rocky planets located in the habitable zone of their host star pose potential for harboring life As of now, we have detected 1,958 different exoplanets These planets live in 1241 planetary systems: there are 488 multiple-planet RV or RV+T systems so far. 249 have two; 82 have 3, 29 have four, 11 have five, 3 have six, 2 have seven. Seen so far, that is.

24 Tidal Forces Take a planet orbiting a star (or a moon orbiting a planet). Gravity gets weaker the further you move from the star. What does that mean for our orbiting planet? Star R1 R2 Planet The gravitational pull felt by the planet due to the star will be greater on its near side than on its far side. The differences in these forces felt by the planet are called tides. Planet Stretch and squish greatly exaggerated Toward orbital Gravitational forces partner (star, planet Forces from the or moon) planet s viewpoint

25 Tidal Forces (cont) The stretching/compressing of the planet due to the tidal forces takes work Not all of the energy is restored when the forces decrease due to friction and viscosity within the planet This drains rotational and orbital energy causing the planet (or moon) to rotate more slowly AND to drift to a larger orbital distance (also circularizes the orbit) After a while (millions of years), this results in a special relation (equilibrium) between the orbital and rotation period called tidal locking. Orbital period = Rotational period The Pluto-Charon System (Wikipedia)

26 Tidally Locked Worlds If a planet is tidally locked, one side will forever face the sun and the other face away This could create extreme desert conditions on the day side, and oppositely extreme icy conditions on the night side (with a potential goldilocks zone in between in which liquid water could exist) Tidal forces can also heat the interior of the planet (like squeezing a tennis ball) This is particularly interesting when considering moons of gas giant planets Tidal heating could warm an icy moon s interior to the point where an entire subsurface ocean of liquid water could exist Artist s conception of a tidally locked planet (space.com) This is exactly the case with Jupiter s moon Europa and Saturn s moon Enceladus Thus allowing life to exist outside the habitable zone

27 Don t forget! Review the practice exam Go over old homework problems Look over the lecture slides Start the exam when you re ready, but give yourself time to finish (6:00pm is the cutoff, if you start at 5:55, you will have 5 minutes to take your exam) Good Luck!!!