Astronomy Study Guide Exam 2
Astronomy Study Guide Exam 2 ASTRON 0089
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Astronomy Section 2 3215 1122 AM Chapter 9 Probing the Dynamic Sun The Sun Diameter 14 X 10A6 km 109X diameter of earth Mass 20 X 10A30 kg 332000X mass of earth Surface temperature 5800 K Core temperature 15000000 K Solar Energy The Sun s spectrum is close to that of an idealized blackbody with a temperature of 5800 K It shines because hundreds of millions of tons of hydrogen are squeezed together and converted to helium every second at its core Luminosity total energy output per second 39 X 1026 jouless 39 X 1026 Watts detonation of 100 billion lmegaton nuclear bombs per sec The Source of the Sun s Energy nuclear reactions occur in the core of the Sun The reaction is called the protonproton chain Hydrogen is converted to Helium Energy is generated by thermonuclear fusion 0 Can take place only at extremely high temperatures because in extreme heat and pressure at the Sun s center positively charged hydrogen nuclei are moving so fast they can overcome their electric repulsion and touch one another and combine creating larger nuclei and emitting energy The temperatures and pressures deep Within the core of the Sun are so intense that hydrogen nuclei can combine fuse to produce helium nuclei Visible light is emitted with the most intensity Assume that the sun is made of only hydrogen for arguments sake and that all of it will be eventually converted in to energy 07 is converted into energy 07 of the suns mass is making energy Energy 007 X 2 X1 0A30 X 3X10A8A2kg msA2 13 X 10 quot44 Joules o The entire sun can give you this much energy Time 13X10A44j0u165 4X10A26jouless 3X10 A17 or 10A10 years seconds to burn out NEED TO KNOW ORDER OF THE LAYERS OF THE SUN Transition zone outside Chromosphere Photosphere Convection zone Radiation zone Core inside Converting Hydrogen to Helium 4 H 9 He energy The particles are 0 Positrons e positively charged electrons antielectrons I Positrons instantly encounter electrons and they annihilate to create a gamma ray 0 Neutrinos v weakly interacting neutral particles that travel at nearly the speed of light and have nearly no mass 0 Gamma ray photons y energy 0 Deuterium HA2 an isotope of Hydrogen with l proton and l neutron o Helium3 HeA3 an isotope of Helium with 2 protons and l neutron The sum of all the masses created in this reaction is less than the mass of the particles that were used up Mass of 4 hydrogen nuclei 66943 X 10A27 kg Mass of l Helium4 nucleus 66466 X 10A27 kg The lost mass 07 has been converted to energy Einstein s massenergy equation based on his theory of special relativity o E mcquot2 o E amount of energy into which mass can be converted Joules o m quantity of mass kg 0 c speed of light 3 X 10A8 ms A small amount of matter can release an incredible amount of energy because the speed of light squared is huge This is the energy that powers the Sun and that eventually radiates out into space The protons overcome their mutual repulsive force bc the temperature in the core of the Sun is very high 15 million K o This means that they move very fast in an extremely dense environment 0 They overcome their mutual repulsion stick together and fuse The energy generated as gamma rays leaves the Sun as visible light but it takes a few hundred thousand years to make its way to the Sun s surface 0 This is bc it is absorbed and reemitted many times as it bounces off matter in the Sun s interior However neutrinos generated in the core travel cleanly through the Sun bc they have no charge and interact very weakly with matter 0 they leave near the speed of light and escape into space a few seconds after being created 10A38 neutrinos are created in the Sun each second Every second about 10A14 neutrinos pass through each square meter of the Earth More Sun Stats Composition 912 H and 87 He by number of atoms 710 H and 271 He by mass plus tiny amounts of 65 other identi ed elements Hydrostatic and Thermal Equilibrium Equilibrium is maintained by a balance among three forces 1 The downward pressure of the layers of solar material 2 The upward pressure generated by hot gases 3 The slab s weight gravitational pull it feels from the rest of the sun The Sun is in hydrostatic equilibrium which means that inward gravitational force is balanced by outward radiation pressure Thermal Equilibrium The temperature structure inside the Sun remains constant in time All the energy generated in the core of the Sun must be transported to the surface from where it s radiated into space Energy increases 9 temperature increases 9 pressure increases 9 surface eXpands 9 more energy radiated out Transporting energy through the Sun The core of the Sun is where its energy is generated This energy is transported by radiation photons in the Radioactive Zone This energy is transported by convection mass motions of gas in the Convective Zone Modeling the Sun To construct a model of a star like the Sun we express the ideas of hydrostatic equilibrium thermal equilibrium and energy transport as a set of equations We then calculate conditions layer by layer in toward the star s center The result is a model of how temperature pressure and density increase with increasing depth below the star s surface Probing the Sun s Interior Heliosiesmology a powerful technique to infer what s going on beneath the Sun s surface involves measuring vibrations of the Sun as a whole The Solar Atmosphere The Photosphere o The visible surface of the Sun that we re most familiar with 0 Since the Sun is a ball of gas this isn t a solid surface it s a layer about 400 km thick 0 The photosphere is the layer in the solar atmosphere from which the Sun s visible light is emitted o The Sun appears darker around its limb or edge bc we see the upper photosphere which is relatively cool and glows less brightly o The dark sunspots are also relatively cool regions Granulation Highresolution photographs of the Sun s surface reveal a blotchy pattern called granulation Granules are convection cells about 1000 km 600 mi wide in the Sun s photosphere Rising hot gas produces bright granules Cooler gas sinks downward along the boundaries between granules this gas glows less brightly giving the boundaries their dark appearance This convective motion transports heat from the Sun s interior outward to the solar atmosphere Supergranules and LargeScale Convection Supergranules display little contrast between their center and edges so they re hard to observe in ordinary images In a falsecolor Doppler image light from gas that s approaching us rising is shifted toward shorter wavelengths whereas light from receding gas descending is shifted toward longer wavelengths Observing Sunspots A sunspot is a region in the photosphere where the temperature is relatively low which makes it appear darker than its surroundings Although sunspots vary greatly in size typical ones measure few tens of thousands of kilometers across comparable to the diameter of the Earth Sunspots are not permanent features of the photosphere but last between a few hours and months Sunspots are associated with magnetic elds 0 Outside the sunspot the magnetic field is low 4 0 Within the sunspot the magnetic eld is strong The Sunspot Cycle The average number of sunspots on the Sun isn t constant but varies in a predictable sunspot cycle Sunspots appear and disappear with time their numbers and distribution across the face of the Sun also change in regular fashion The average number of spots reaches a maX every 11 years then falls off to almost 0 before the cycle begins again the Sun s magnetic eld changes polarity every 11 years What causes sunspots The Sun rotates faster at its equator than its poles differential rotation 25 rotation period at the equator and 36 days at the poles The strong magnetic eld gets distorted as it wraps around the equator The N S magnetic eld eventually becomes an EW f1eld Convection causes the magnetized gas to leave the surface and loop through the atmosphere taking the magnetic eld with it it creates a sunspot pair Sunspots come and go typically in a few days They re linked by pairs of magnetic eld lines Magnetic Arches Plasma charged particles tends to follow the Sun s magnetic eld with streamers of electrically charged particles moving along each eld line When the magnetic elds of 2 arches come into proximity they can rearrange and combine The tremendous amount of energy stored in the magnetic eld is then released into the solar atmosphere Prominences Magnetic elds can also push upward from the Sun s interior compressing and heating a portion of the chromosphere that appears as bright arching columns of gas called prominences These can extend for tens of thousands of kilometers above the photosphere Some prominences last for only a few hours while others persist for many months The most energetic prominences break free of the magnetic elds that confmed them and burst into space The Sun s Chromosphere During a total solar eclipse the Sun s glowing chromosphere can be seen around the edge of the moon It appears pinkish because its hot gases emit light at certain wavelengths principally the Ha emission of hydrogen at a red wavelength of 6563 nm Spicules are jets of chromospheric gas that surge upward into the Sun s outer atmosphere The temperature increases as you go higher in the chromosphere The Corona The outer layer of the Sun s atmosphere extending out to a distance of several million kilometers Only about onemillionth as bright as the photosphere and can be viewed only when the light from the photosphere is blocked out Looks like numerous streamers extending in different directions far above the solar surface changing over days and weeks Has temperatures far greater than the temperatures in the chromosphere Temperatures in the Sun s Upper Atmosphere The corona is actually not very hot containing very little thermal energy The corona is nearly a vacuum but the atoms there are moving at very high speeds But because there are so few atoms in the corona the total amount of energy in these moving atoms measure of how hot the gas is is rather low Air 10A25 particles corona 10All photosphere 10A23 Corona is much hotter than layers below it must have a heat source It s not clear what heats the corona but apparently the magnetic eld of the Sun acts like a pump that increases the speed of particles in the corona Solar Wind Solar wind is the out ow of coronal gases traveling at 1 million kilometers per hour Each second the Sun ejects about a million tons of material into the solar wind composed almost entirely of electrons and nuclei of hydrogen and helium The aurorae seen at far northern or southern latitudes on Earth are produced when electrons and ions from the solar wind enter our upper atmosphere Solar wind escapes Sun mostly through coronal holes which can be seen in Xray images The Sun is evaporating The wind carries away about 2 million tons of solar matter every second But less than 01 of the Sun s mass has been lost this way since the Solar System formed 48 billion years ago Solar corona changes along with sunspot cycle it s much larger and more irregular at sunspot peak Coronal Mass Ejections Huge blasts of high energy particles are followed by vast amounts of solar plasma traveling outward at hundreds of kms In a day or two they reach the orbit of the Earth The Sun s magnetic elds and releases of plasma directly affect Earth and the rest of the solar system Solar wind shapes the Earth s magnetosphere and magnetic storms can disrupt communications and navigational equipment damage satellites and even cause blackouts Producing the Aurorae Some of the high energy particles can get into the Earth s magnetosphere which otherwise protects us They impact our upper atmosphere and it glows The Earth s magnetic field causes particles to stream preferentially down the magnetic poles so normally you have to be at high latitudes to see them The Aurorae Some solar wind particles are able to leak through Earth s magnetic field at its weaker points and cascade down into the Earth s upper atmosphere to where Earth s magnetic field connects with the Earth near the planet s north and south poles As these highspeed charged particles collide with atoms in the upper atmosphere they excite the atoms to high energy levels The atoms then emit visible light as they drop down to their ground states like the excited gas atoms in a neon light Chapter 10 Observing Properties of Distant Stars Fundamental Information about Stars Distance how do we determine distances to stars Brightness how do we measure stellar brightness Temperature Wien s Law Size how do we know how large stars are Mass how massive are they How do we know Chemical composition spectroscopy Age how old are stars How do we know Vast Emptiness of Space Nearest star to the sun Proxima Centauri which is a member of the 3star system Alpha Centauri complex is 43 ly away Scale of distances 0 Sun is a marble earth is a grain of sand orbiting l m away 0 Nearest star is another marble 270 km away 0 Solar system extends about 50 m from Sun rest of space to nearest star is basically empty Parallax Parallax uses the change in an object s apparent position compared to the background Imagine looking at some nearby object tree against a distant background mountains when you move from one location to another the nearby object appears to shift with respect to the distant background scenery Stellar Parallax As the Earth orbits the Sun a nearby star appears to shift its position against the background of distant stars The parallax of the star is equal to the angular radius of the Earth s orbit as seen from the star The closer the star is to us the greater the parallax angle Parallax and the distances to stars d 1 p d distance to a star in parsecs p parallax angle of that star in arcseconds This simple relationship between parallax and distance reveals that the closest stars have the greatest parallax Astronomers often describe distances in parsecs rather than lightyears because they use the observed parallax to measure distance 1 arcsecond l60th arcminute A Parsec The parsec a unit of length commonly used by astronomers is equal to 326 light years 0 1 pc 326 ly The parsec is defined as the distance at which 1 AU perpendicular to the observer s line of sight subtends an angle of l arcsec o l arsec l l60 13600 degrees 1 parsec 206265 AU Luminosity and Apparent Brightness Luminosity absolute brightness measure of the total energy radiated by a star 0 Units energy per second Apparent brightness how bright a star appears when viewed from earth Depends on luminosity and distance to the star The farther away you are from a light source the dimmer it looks How does light spread outThe Inverse Square Law of Radiation The amount of light received gets diluted by the square of the distance from the source The same amount of energy is spread out over a larger and larger sphere as light travels away from a star 0 The area of a sphere is 47TI39A2 Apparent brightness b intrinsic brightness 47tDA2 L 47tDA2 D distance between star and observer L luminosity intrinsic brightness energy radiated per second b apparentobserved brightness If we know a star s apparent brightness and its distance we can calculate its luminosity Brighter doesn t always mean more luminous Two stars that appear equally bright to us might be a closer dimmer star and a farther brighter one greater distance can counteract greater luminosity Star Colors and Temperature Red stars are cooler blue stars are hotter A star s color depends on its surface temperature The intensity of light from relatively cool objects peaks at long wavelengths making the star look red A hot star s intensity curve peaks at short wavelengths making it look blue Classifying Stars Spectral Classes The atmospheres of stars produce absorption line spectra depending on the presence of different elements These spectra are diverse allowing us to classify stars into spectral classes 0 OBAFGKM 0 Oh Be A Fine Guy Kiss Me Spectral types reveal temperatures Based on the structure of atoms we understand that OBAFGKM spectral sequence is actually a sequence in temperature The hottest stars are 0 stars their absorption lines can occur only if these stars have surface temperatures above 25000 K The coolest stars are M stars the spectral features of M stars are consistent with stellar surface temperatures of about 3000 K In other words the sequence OBAFGKM is also a temperature sequence from hottest to coldest Surface temperature affects stellar spectra For hydrogen lines to be prominent in a star s spectrum the star must be hot enough to excite the electrons out of the ground state but not so hot that all the hydrogen atoms become ionized A stellar surface temperature of about 9000 K produces the strongest hydrogen lines this is the case with A0 and A5 stars Every other type of atom or molecule also has a characteristic temperature range in which it produces prominent absorption lines in the observable part of the spectrum Spectral Classes for Brown Dwarfs Brown dwarfs too small to sustain thermonuclear reactions but too big to be considered planets Observing in the infrared 2 new spectral classes are added for brown dwarfs L and T o The modern classes are OBAFGKMLT Concept Check 1 What quantity dictates what spectra class you give to a star Temperature 2 Is a spectral class F2 star more similar to an Aspectral class star or a Gspectral class star A 3 Which brown dwarf is hotter an L or a T L Stellar Sizes Some stars are close enough and big enough that we can directly measure their sizes For example Betelgeuse is hundreds of times larger than the sun Stellar radii vary widely Relationship between a star s luminosity radius and surface temperature L 4nrA2 6TA4 L star s luminosity Watts R star s radius meters 6 StefanBoltzmann constant T star s surface temperature K StefanBoltzmann s Law energy radiated sec per mquot2 E oTA4 If the distance is known then the intrinsic luminosity is also known Inverse Square Law of Radiation The radius size of the star can then be calculated 10 If R radius of star then total surface area of star 47IRA2 o L 47TRA26TA4 energy radiated per sec J oules s T can be determined from the spectrum of the star This method is an indirect determination or R If 2 stars have the same luminosity but A has a twice the surface temperature A s radius is 1A as large as B What does this mean A cool star can be bright if it has a very large radius A very hot star can be dim if it has a very small radius Binary Stars Orbit Each Other A seesaw balances if the fulcrum is at the center of mass of the 2 children The members of a binary star system orbit around the center of mass of the 2 stars Although their elliptical orbits cross each other the 2 stars are always on opposite side of the center of mass and thus never collide The Concept of Center of Mass 0 M1R1 M2R2 R1 distance of mass Ml from center of mass R2 distance of mass M2 from center of mass Binary Stars The middle star of the handle is binary Most stars are members of multiple star systems the majority of stars are found in binary pairs Visual binaries can be measured directly Ex Mizar A Kruger 60 Sirius is a binary o Sirius A has a smaller orbit but is larger 0 Sirius B has a larger orbit but is smaller and further away from center of mass The more massive object makes a smaller circle and vice versa Eclipsing binary star one star comes in front of the other 0 Can be used to measure size 0 If we know how far away the stars are we can calculate how big they are 0 Shape and timing of eclipse gives sizes of stars Binary stars are very useful for determining the masses of stars Orbital parameters such as period and size of the orbit and velocities of the 2 stars can sometimes be observed We can then use Newton s theory of gravity to calculate stellar masses 11 F Gmlm2 rA2 Binary Star Systems Reveal the Masses of Stars There s no practical direct way to measure the mass of an isolated star Binary star binaries are pairs of stars that orbit each other Their mutual gravitational attraction causes their orbital motions Newton s law of gravity can be used to get 0 M1 M2 masses of2 stars 0 a semimajor axis in AU 0 P orbital period in years 0 Kepler s Third Law of Planetary Motion I M1M2aA3PA2 Spectroscopic binaries can be measured using their Doppler shifts Motion away from observer causes redshift Motion toward observer causes blueshift The HertzsprungRussell HR Diagram This is a plot of stellar luminosity vs surface temperature Some wellknown stars are plotted here Temperature changes from low to high 3000 K to 30000 K Each star is represented by a dot The position of each dot corresponds to its luminosity and temperature The vertical position represents the star s luminosity The horizontal position represents the star s surface temperature Once many stars are plotted on an HR Diagram a pattern begins to form 9 Diagram of closest 80 stars Dark curve on the graph is called the main sequence where most stars are located The white dwarf region is where very hot stars are located but they aren t very luminous since they re small 9 Diagram of 100 brightest stars All more luminous than the sun Two new categories appear red giants and blue giants The brightest stars in the sky appear bright because of their enormous luminosities not their proximity 9 Diagram of 20000 stars Main sequence and red giant region are clear About 90 of stars lie on main sequence 9 re red giants and 1 are white dwarfs Stellar Masses 12 Mass is the main determinant of where a star will be on the Main Sequence This pie chart shows the distribution of stellar masses The more massive stars are much rarer than the least massive Mass is correlated with radius and is very strongly correlated with luminosity for stars that lie on the Main Sequence The MassLuminosity Relation The greater the mass of a mainsequence star the greater its luminosity its surface temperature and its radius A main sequence star of mass 10 M has roughly 3000 times the Sun s luminosity 3000L one with 01 M has a luminosity of only about 0001 L Mass is also related to stellar lifetime The more mass the shorter its lifetime Stellar lifetime l stellar mass3 So the most massive stars have the shortest lifetimes they have a lot of fuel but burn it at a very rapid pace On the other hand small red dwarfs burn their fuel extremely slowly and can have lifetimes of a trillion years or more Chapter 11 The Interstellar Medium Dust clouds dark regions in the Milky Way that block light from the stars beyond M8 The Lagoon Nebula M20 The Trif1d Nebula M17 The Omega Nebula M16 The Eagle Nebula Nebula general term for fuzzy objects in the sky 0 Dark Nebula dust cloud 0 Emission Nebula glows due to hot stars near them I HII Regions different parts of nebula red green and blue 0 Reflection Nebula blue due to light scattering by dust I Lower concentrations of dust than dark nebulae Interstellar medium consists of gas and dust Gas atoms and small molecules mostly hydrogen and helium cold less than 100K Dust sootsmoke larger clumps of particles 0 Absorbs light extinction and reddens light that gets through reddening o Reddening can interfere with blackbody temp measurement but spectral lines don t shift dust only changes color ISM gas is very diffuse 13 Dust is many light years deep which is why we can t see through it Interstellar dust grains complex in shape 0 Grow linearly and become rodlike on small scales but become tangled and twisted in complex ways on larger scales How Nebulae Work Strong interaction bt nebula and stars within it photoevaporation Light from hot stars scatters to dust cloud which appears blue to you because of the re ected scatter light star would appear red Emission nebulae are made of hot thin gas which exhibits distinct emission lines Some emission lines come from forbidden transitions they re so rare that under standard lab conditions they re never seen Forbidden transition in doubly ionized oxygen 0111 is responsible for green in Orion nebula Dark Dust Clouds Average temperature of dark dust clouds is a few tens of Kelvins Absorb visible light top and emit at radio wavelengths bottom Emission os from the CO molecule Ex Ophiuchus dust cloud irregular shape with streamers o Horsehead Nebula famous 21Centimeter Radiation Interstellar gas emits lowenergy radiation due to a transition in the hydrogen atom The emitted photon has a wavelength of 21 cm which is in the radio portion of the electromagnetic spectrum Proton amp electron spin parallel to each other and then ip to a more stable state antiparallel spins emitting a photon which has that 21 cm wavelength Interstellar Molecules The densest gas clouds are also very cold around 20 K These clouds tend to contain molecules rather than atoms Transitions between rotation states faster to slow of a molecule emit radio frequency photons Radio waves aren t absorbed much so molecular gas clouds can be detected even though there may be other gas and dust clouds in the way These clouds are mostly molecular hydrogen which doesn t emit in the radio portion of the spectrum but other molecules do Molecular cloud complexes are enormous around 50 pc across and have enough material to make a million Suns 14 0 There are about 1000 known in the Milky Way Star Formation and Evolution The Formation of Stars Like the Sun Stage 1 Interstellar clouds start to contract triggered by shock or pressure wave from nearby star or supernova As it contracts cloud fragments into small pieces eventually forming many tens or hundreds of individual stars Stars form in clusters Stage 2 Individual cloud fragments begin to collapse Once density is high enough no further fragmentation Stage 3 Interior of the fragment has begun heating and is about 10000 K Stage 4 Core of cloud is now a protostar and makes its rst appearance on HR Diagram 0 It now has an L and T that are within range of this plot The dusty interstellar cloud condenses under selfgravity Gravitational energy is converted to heat and light energy 9 this is a protostar Since gravity is strongest at the center the center condenses faster and heats up faster Central density and pressure go up Energy source gravity CORE TEMP OF SUN 15 MILLION K Planetary formation has begun but the protostar is not in equilibrium all heating comes from the gravitational collapse Stage 5 protostar s luminosity decreases even as its temperature rises bc it s becoming more compact Stage 6 core reaches 10 million K and nuclear fusion begins 9 protostar has become a star Object that can produce its on thermonuclear fusion Stage 7 star continues to contract and increase in temperature until it s in equilibrium Star has reached the Main Sequence and will remain there as long as it has hydrogen to fuse The Zero Age Main Sequence ZAMS The star is in equilibrium The inward pull of gravity is balanced by the gas pressure T core 15 million K T surface 5800 K 15 R 1 R Energy Source pp chain Observations of Cloud Fragments and Protostars Emission nebulae are heated by the formation of stars nearby Protostars have very strong winds jets which clear out an area around the star roughly the size of the solar system 9 these are T Tauri stars Proplyds protoplanetary disks Important stars don t move along the Main Sequence once they reach it they are in equilibrium and don t move until their fuel begins to run out When a star moves in the HR Diagram it means its L andor T change so that it now is at a different spot on the diagram Since L and T change smoothly it appears as if the star moves as it evolves Stars of other masses Stars less massive than the Sun have similar path shapes in the HR Diagram but they wind up in different places on the Main Sequence Pre Main Sequence evolution is Hayashi Track Some fragments are too small for fusion ever to begin they gradually cool off A protostar must have 008 M 80x the mass of Jupiter to become dense and hot enough that fusion can begin If the mass of the failed star is about 12 Jupiter masses or more it s luminous when first formed and called a brown dwarf Star Clusters Because a single interstellar cloud can produce many stars of the same age and composition star clusters are an excellent way to study the effect of mass on stellar evolution The stars are also all at the same distance from us NGC 3603 newborn cluster of a hot young blue Type 0 and B stars an OB Association The Pleiades is a young star cluster an example of an open cluster 0 Most stars in this are on the Main Sequence 0 The brightest stars are the ones with luminosities over 1000 L 0 They are blueish and white B and A stars which we see Globular cluster absence of massive Main Sequence stars heavily populated Red Giant region Relatively old contain stars of all masses Cluster HR Diagrams give the age of the cluster Inhabit the halo of our galaxy Stellar Evolution 16 Eventually as hydrogen in the core is consumed the star begins to leave the Main Sequence Its evolution from then on depends on the mass of the star 0 Lowmass stars go quietly o Highmass stars go out with a bang Evolution of a Sunlike Star Even while on Main Sequence the composition of a star s core is changing Helium begins to accumulate over time due to Hydrogen fusion The sun is a little bigger today than when it first formed Most of the core is Helium The core shrinks gets hotter amp denser the rate of Hydrogen fusion increases luminosity increases radius increases The Sun is now 40 more luminous and 6 larger Stage 8 As the fuel in the core is used up less radiation pressure can be produced and the core contracts remember constant tug of war bt gravity and radiation pressure When the fuel is used up nothing holds up the star so the core begins to collapse due to gravity Hydrogen in a shell around the core heats up and begins to fuse to form Helium It s now burning hotter than before and radiation pressure pushes the outer layers farther out Stage 9 The RedGiant Branch As the core shrinks it heats up the shell of Hydrogen outside it burns vigorously and generates much heat This makes the outer layers of the star expand and cool The sunlike star is now a red giant Despite its cooler temp its luminosity increases enormously due to its large size Stage 10 Helium Fusion Once the core temp has risen to 100000000 K the helium in the core starts to fuse through the triplealpha process 3 helium nuclei fuse to form carbon 0 Beryllium is produced very brie y 4He 4He 9 8Be energy 8Be 4He 9 12C energy The 8Be nucleus is highly unstable and will break up into 2 helium nuclei in about 10A12 seconds unless another helium nucleus fuses with it first This is why temperatures and densities are necessary 17 The Helium Flash Stage 11 Electron Degeneracyz two electrons can t be in the same quantum state they can t be squeezed together beyond a certain point So as the core collapses due to gravity the electrons reach this point of degeneracy and hold up the star When helium starts to burn in the core the temperature rises but the gas can t expand to release pressure in response to the rise in temperature because electron degeneracy has made the core rigid So the temperature rises even faster it s like putting a tight lid on a pot of boiling water Helium begins to fuse extremely rapidly nally the temperature gets high enough to break the electron hold and let the core expand again Within hours the enormous energy output is over and the star once again reaches equilibrium Back to the giant branch As the helium in the core fuses to carbon the core becomes hotter and hotter and the helium burns faster and faster The star is now similar to its condition just as it left the Main Sequence except now there are two shells The star has become a red giant for the second time Such stars like the Sun never become hot enough for fusion past carbon to take place There is no more outward radiation pressure due to fusion being generated in the core so the core continues to contract The outer layers become unstable and are eventually ejected Over 1000 years the star loses 40 of its mass and exposes the 100000 K core Contraction of the core stops when electron degeneracy kicks in again the electrons hold up the star But this time there is no helium burning taking place that can increase the temperature and unlock the electrons The core is now called a white dwarf the star is dead It doesn t produce any energy anymore it cools off The Death of a LowMass Star Stage 12 The ejected envelope expands into interstellar space forming a planetary nebula why we see a ring The star now has two parts 0 A small very dense carbon core 18 0 An envelope about the size of our solar system planetary nebula As the dead core of the star cools the nebula continues to expand and dissipates into the surroundings Stages 13 amp 14 White and Black Dwarfs Once the nebula has gone the remaining core is extremely dense and hot but quite small It s luminous only due to its high temperature White dwarf shell burning dies out leaving behind a CO core that contracts due to the force of gravity 0 The contraction is stopped by electron degeneracy pressure 9 it s now a white dwarf o It weighs a ton a teaspoon Density 1 million gcmA3 R R earth T few 10000 K Energy source cooling off The Hubble Space Telescope has detected white dwarf stars in globular clusters As the white dwarf cools its size doesn t change signi cantly it simply gets dimmer and dimmer until it ceases to glow Stars of the same age but different mass can appear as the whole cluster ages After 10 million years the most massive stars have already left the Main Sequence while many haven t even reached it yet After 100 million years a distinct mainsequence turnoff begins to develop this shows the highestmass stars that are still on the Main Sequence After 1 billion years the mainsequence turnoff is much clearer After 10 billion years a number of features are evident o The redgiant subgiant asymptotic giant and horizontal branches are all clearly populated White dwarfs also appear indicating that solarmass stars are in their last phases Hyades Cluster young mainsequence turnoff indicates an age of 600 million years 47 Tucanae about 1012 billion years old 19 20 3215 1122 AM 21 3215 1122 AM
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