08 Supernovae Remnants I

Transcript

High Energy Astrophysics Supernovae and their Remnants 1/2 Giampaolo Pisano Jodrell Bank Centre for Astrophysics - University of Manchester [email protected] March 2012 Supernovae and their remnants - Introduction - Evolution of low mass stars - White dwarfs - Evolution of high mass stars - Supernovae Type II - Supernovae Type Ia - Supernova Remnants (SNR) References: - Longair - Vol 2 - Par. 13.3, 14.2, 14.3 & 15.3 - Freedmann - Kaufmann, Universe - Chap. 20 - Rosswog & Bruggen - Chap.4 Supernovae 1/2 - SN are catastrophic explosions of stars:
Within a few weeks a SN can produce as much light as its progenitor star during its entire life - The classification of SN is complex, anyway the two major types are: Presence of H lines in spectra SN Type II Light curves quite diverse Mechanism: core collapse explosions Absence H lines SN Type Ia Light curves all similar Mechanism: thermonuclear explosions Note: there are also SN Type Ib and Ic that are core-collapse SN Supernovae 2/2 - Occurence of Supernovae: SN within our Galaxy - Occur every 30-50 yr (many obscured by interstellar gas and dust) - Known SN explosions in: 1006, 1054, 1181, 1572, 1604 & 1680 - During existence Milky Way
~108 SN have exploded SN in the Universe - SN recently observed with known progenitor:
SN1987A, SN1993J, SN2005CS - Approximately one SN explosion every second Extragalactic Supernovae 1/2 - SN1994D Hubble Space Telescope NASA/ESA - SN explosion in galaxy NGC 4526 - Supernovae have peak luminosities of the order of 109 Lʘ:
these are comparable to galaxy luminosities Extragalactic Supernovae 2/2 - Supernovae can be observed in distant galaxies
SN important to measure large distances in the Universe Supernovae in HEA - These events are important for HEA because: - Birth events of neutron stars and stellar black holes - Powerful source of heating for ambient interstellar gas - Intensive X-ray sources (Bremsstrahlung of hot gas) - Radio sources (Synchrotron of electrons in SNR magnetic fields) - Possible sources of high energy particles - Origin of the most heavy elements in nature Supernovae and their remnants - Introduction - Evolution of low mass stars - White dwarfs - Evolution of high mass stars - Supernovae Type II - Supernovae Type Ia - Supernova Remnants (SNR) References: - Longair - Vol 2 - Par. 13.3, 14.2, 14.3 & 15.3 - Freedmann - Kaufmann, Universe - Chap. 20 - Rosswog & Bruggen - Chap. 4 - Bradt - Par. 2.6, 3.6 & 4.3 The Hertzsprung-Russel Diagram 1/2 Wikipedia The Hertzsprung-Russel Diagram 2/2 http://astro.berkeley.edu H-R diagrams of two clusters look different: - the open cluster M67
a young cluster - globular cluster M4
an old cluster
Turnoff point in the main sequence used to calculate the age of the cluster Evolution of stars in the HR diagram - H-R diagram M >8Mʘ Fusion Ne, O, Si
.. Ni, Fe 4Mʘ1.5 Mʘ
central temperature higher
CNO cycle Energy production 1/3: p-p chain - Proton-proton chain (pp1) Stars with: M < 1.5M Ο. T < 2×107 K First reaction: production of 2H and 3He Principal source of solar neutrinos Wikipedia There are other two reaction chains: pp2 and pp3 Energy production 2/3: CNO cycle - Carbon-Nitrogen-Oxygen cycle Stars with: M > 1.5M Ο. T > 2×107 K Four protons fuse using carbon, nitrogen and oxygen isotopes as catalysts to produce one alpha particle, two positrons and two electron neutrinos. Wikipedia Energy production 3/3: Example - Main sequence lifetime of the Sun - The total energy liberated in the conversion H
He is: ∆E = 0.007 mc 2 - At the ‘turnoff point’ 10% of the Sun mass will be converted in He: (generally known as Chandrasekhar-Schoenberg mass limit:10-15%) ∆EΟ. = 7 ×10 −4 M Ο. c 2 - Assuming the Sun luminosity roughly constant in the main sequence:  LΟ. = 3.9 ×10 26W ∆EΟ. 7 ×10 −4 × 2 ×1030 × (3 ×108 ) 2 → tΟ. ≈ =  26 30 3 . 9 × 10 L . M = 2 × 10 kg Ο .  Ο tΟ. ≈ 1010 years - In the main sequence, for star masses Mʘ < M < 10 Mʘ : −3 4 26  M   Watts L ~ 4 ×10   M Ο.  tlife  M   Gyr ~ 9 ×   M Ο.  Evolution of stars with mass M < 8Mʘ 2/7 M ≅ M Ο. Fusing H shell He core H,He Core contraction Envelope expansion Ignition shell H Contraction He core - When H
He core fusion ceases
core shrinks becoming degenerate: - Heating surrounding H
Shell H fusion starts - New energy causes outer envelope expansion and cooling
Increase in luminosity
the star enters the ‘Red giant phase’ Evolution of stars with mass M < 8Mʘ 3/7 M ≅ M Ο. He core burning He flash Fusing H shell H,He Fusing He Wikipedia - He-rich core mass increases, shrinks and heats up until helium core fusion begins explosively in the core
Helium flash - Core expands cooling down
shell H release less energy - Contraction outer layers
increase surface temperature
Horizontal branch Evolution of stars with mass M < 8Mʘ 4/7 M ≅ M Ο. Ignition shell He Core contraction Envelope expansion Fusing H shell Fusing He shell H,He C,O Core - All He in the core is converted to Carbon and Oxygen - Core contracts until becomes degenerate
heat
Shell helium fusion - Star enters second red-giant phase
outer envelope expansion and cooling Evolution of stars with mass M < 8Mʘ 5/7 M ≅ M Ο. - At the end of the giant phase: - The interaction between the fusing shells generates thermal pulses - The star undergoes bursts in luminosity ejecting the outer layers - Formation of
‘Planetary nebulae’ with exposed ‘hot core’ in the centre Planetary Nebulae 1/2 M57 – Ring Nebula - Nebula is located at 2300 light-years from Earth - Expulsion materials happened approximately 1600 years ago - It is expanding at a rate of approximately 20-30 km/s - Mass approximately 1.2 solar masses - Central white dwarf illuminating the ejected material Planetary Nebulae 2/2 Hubble Space Telescope - The rate of expansion of the gases can be anisotropic - Ejection of gases in different ways and at different stages
More complex shapes Evolution of stars with mass M < 8Mʘ 6/7 M ≅ M Ο. - The star ejected up to 60% of its original mass - The small Carbon-Oxygen hot core get exposed
Rapid increase in Temperature Evolution of stars with mass M < 8Mʘ 7/7 M ≅ M Ο. - There are not thermonuclear reactions in the stellar core - The core cools down without collapsing
it maintains the same size - The final object is called
‘White Dwarf’ White dwarf - Sirius Sirius A Sirius B Hubble Space Telescope - Sirius is the brightest star in the sky - It is actually a binary star: - Sirius A: a main sequence star with surface T=10500 K - Sirius B: a white dwarf with surface T=25200 K White dwarfs 1/12 - White dwarf typical characteristics: - Maximum mass: - Typical size : M WD ≤ 1.46 M Ο. - Chandrasekhar Mass RWD = 5000 − 10000 km - The typical density is: ρWD = ( Earth size ! ) M 1.46 M Ο. = = 6 ×109 kg / m 3 6 3 V 4 3 π (5 ×10 ) At such high densities quantum mechanical effects become important The matter becomes degenerate (not ideal gas anymore) White dwarfs 2/12 - Degenerate plasma pressure - In a white dwarf there is no internal heat source but: Very high density Heisenberg’s uncertainty principle ∆x ∆p ≈ h + Pauli’s exclusion principle Fermions The star is held up by degeneracy pressure Bosons White dwarfs 3/12 - Origin of the degenerate gas pressure (simple qualitative description) - Imagine an electron-proton plasma inside a box in thermal equilibrium at T T T p e- ∆xe ∆xp ∆xe >> ∆x p - In thermal equilibrium: 1 3 m v 2 = kT → m 2 v 2 = 3mkT → p ≈ 3mkT 2 2 pe me 1 ≈ ≈ pp m p 43 - Assuming the average momentum of the order of its uncertainty: h ∆p ≈ p → ∆x ≈ ∆p h ∆x ≈ 3mkT mp ∆xe ≈ ∆x p me
QM effects on electrons become important at larger interparticle distance White dwarfs 4/12 - Quantum mechanical effects become important when: (*) Let’s neglect spin - Inter-particle distance becomes small
Heisenberg principle & - Particle cannot occupy the same state
Pauli principle (*) & (**) T p e- p ∆xe ∆xp Increasing density Electrons become degenerate at much larger interparticle distance ↓ ∆xe e- ↑ ∆pe - If electron spacing ∆xe very small
they have large momenta: ∆p ≈ h / ∆x Electron degenerate pressure dominant (Independent of T) (**) We neglected the possibility to have electrons in the same place with different momenta because ∆p~p White dwarfs 5/12 - Degenerate matter critical density - Most of the mass comes from protons but minimum volume from electrons: mp ∆xe3 ρc ≈ mp ∆x 3 e = mp h  3m kT  ρ c ≈ m p  e2   h  3 (3me kT ) 3 / 2 3/ 2 ∝ T 3/ 2 - Critical density - Density at which degeneracy occurs in the non-relativistic limit Note on white dwarfs - WD constitutes of different chemical elements - WD outermost layer is non-degenerate, anyway it is very thin:
we can assume the WD to be completely degenerate
Electrons provide the degenerate pressure
Protons and nuclei provide the gravitational pull White dwarfs 6/12 - In general, for any chemical composition the exact calculation gives: T  ρ c = 2.38 ×10 −5   µ 3/ 2 µ e5 / 2 kg / m 3 - Critical density where: # nucleons  µ =  e # free electrons    # nucleons µ =  # free particles  Ionised hydrogen plasma: 1  µ =  e 1 =1  1 µ = = 0.5 2  Ionised helium plasma: 4  µ =  e 2=2  4 µ = = 1.3 3  - The density of a material can be written as: ρ ≅ (n + ne ) µ m p = ne µe m p  n : number density of nuclei  ne : number density of electrons White dwarfs 7/12 - 1D Fermi-Dirac electron gas Bradt Fig.3.5 Non-degenerate Maxwell-Boltzmann 1D distribution Degenerate Pressure nearly independent on T Momenta higher than normal gas Small ∆T does not change the pressure
Fermion gas degenerate if (at sufficiently low T) or (with a sufficiently high density) White dwarfs 8/12 - Fermi energy and temperature for a 3D degenerate gas - For a degenerate gas: h h  1/ 3 → ∆ p ∝ ∝ n p ∆ ≈ e  ∆x ∆x  1 ne ≈ 3 → ∆x ≈ 11/ 3 ∆x  ne 1/ 3 It can be shown: - Fermi momentum 2 - The Fermi energy is then:  3n  p F = h e   8π  pF h 2  3ne  EF = =   2me 2me  8π  2/3 - The Fermi temperature is related to the Fermi energy: TF =
A plasma is degenerate if its temperature is: T ≤ TF
The pressure is independent from temperature EF k White dwarfs 9/12 - Degenerate matter in a white dwarf Helium white dwarf  µe = 2  9 3 ρ = 6 × 10 kg / m  WD ρWD 6 ×109 ne ≅ = 1.8 ×1036 m −3 = - 27 µe m p 2 × 1.7 ×10 h 2  3ne  EF =   2me  8π  2/3 ≈ 8.7 ×10 −14 J TF ≈ 6.3 ×109 K - White dwarfs are so called because their emission peaks around: λmax ~ 400nm Blackbody → T = 0.003 / λmax = 7500K TWD << TF T = 7500 K - Typical White Dwarf Surface Temperature
The interior of a white dwarf is degenerate A white dwarf does not collapse because of the electrons Fermi pressure White dwarfs 10/12 - At extremely high pressure the electrons can become relativistic - The correct calculation leads to three different states of matter: - Non-degenerate (normal) - Degenerate but non-relativistic - Degenerate and relativistic
The state depends on the temperature and density White dwarfs 11/12 - Types of Equation of State ( Longair Fig.15.8 ) - Radiation pressure can exceed the gas pressure (some hottest massive stars) - In the Sun the equation of state can always be that of a classical gas - The core of stars in the ‘giant branch’ can become degenerate - A degenerate gas can become a solid
ex: cold white dwarfs - For a relativistic degenerate gas the pressure for a given density is smaller White dwarfs 12/12 - There is a limit on the mass of a white dwarf If density too high Gas becomes relativistic Fermi pressure drops Gravity can exceed Fermi pressure Core collapse ! - The largest mass sustainable by Fermi pressure is: M Chandr = 5.836 µ 2 e M Ο. - Chandrasekhar Mass - For a He white dwarf µe=2 - White dwarfs with heavier elements µe can increase slightly - Anyway, for fully ionised iron: µe = 56 =2 28
~ Same mass limit Forming white dwarf in young planetary nebula - NGC 6543: The Cat's Eye Nebula Redux (Chandra X-ray Obs) Composite image Chandra - HST http://chandra.harvard.edu/ - Optical: the radiation pressure from the hot core pushes the ejected materials outward creating the filamentary structures - X-ray: the central star is surrounded by a cloud of multi-million-degree gas
The star is expected to become a white dwarf star in a few million years Young white dwarfs in nearby galaxy - Hot White Dwarf in Young Star Cluster NGC 1818 (HST) - Over 20000 stars in the young (40 million year old) cluster in the Large Magellanic Cloud (LMC), a satellite galaxy of our Milky Way
Isolated white dwarf within the cluster - The LMC is a site of intense star formation: ideal for studying stellar evolution Old white dwarfs in globular clusters 1/2 - White Dwarf Stars in Globular Cluster M4 (HST) - Globular clusters contain hundreds of thousands of old stars - Right image: small portion of the cluster
0.63 light-years across - HST found 75 white dwarf stars in the area it viewed - 40000 are predicted for the cluster as a whole
WDs as the oldest burned-out stars in our Galaxy: 12-13 billion yrs old
Completely independent reading of the universe age without relying on measurements of the universe expansion Old white dwarfs in globular clusters 2/2 - Faintest Stars in Globular Cluster NGC 6397 (HST) White dwarf Red dwarf star - Among brilliant stars are the faintest stars ever seen in a globular cluster: - the faintest red dwarf stars (26th mag) - the dimmest white dwarfs (28th mag): equivalent to the light produced by a birthday candle on the Moon as seen from Earth White dwarfs binary system - RX J0806.3+1527 (Chandra X-ray Obs) http://chandra.harvard.edu/ - X-ray intensity observed with a period of 321.5 seconds
white dwarfs binary star system with orbit period ~ 5 minutes
distance ~50000 miles, velocities > 106 miles per hour - Gravitational waves should be produced:
Energy loss
stars closer and closer at a rate of ~ 2 feet / day
Orbital period is decreasing by 1.2 ms / year White dwarfs types - Spectral classification http://www.daviddarling.info/ - Lighter stars (M<4Mʘ)
carbon-oxygen white dwarfs - Heavier stars (M>4Mʘ)
neon-oxygen dwarfs - In addition WDs differ in terms of their spectra which are dictated by the elements that dominate their surfaces (ex: DA, DB, DC, DO, DZ, DQ) White dwarfs cool down - How does a white dwarf radiate its energy ? There are not thermonuclear reactions anymore There is not gravitational collapse because of the electrons degeneracy The degenerate electrons move very fast but they do not lose energy The non-degenerate nuclei can radiate energy White dwarfs luminosity due to thermal energy of ions (heat capacity) After ~10 billion years the luminosity drops Black dwarf White dwarf crystallisation - Variable white dwarf - BPM 37093 Wikipedia - As a white dwarf cools, its material should crystallize, starting at the centre - BPM 37093 is a variable white dwarf star thought to be composed of C and O - Observations of star pulsations:
information about its internal structure
estimated that a consistent percentage of the mass had crystallized - Crystallization of this type of WD is thought to be a lattice of carbon and/or oxygen nuclei, surrounded by a Fermi sea of electrons (huge diamond!) Supernovae and their remnants - Introduction - Evolution of low mass stars - White dwarfs - Evolution of high mass stars - Supernovae Type II - Supernovae Type Ia - Supernova Remnants (SNR) References: - Longair - Vol 2 - Par. 13.3, 14.2, 14.3 & 15.3 - Freedmann - Kaufmann, Universe - Chap 20 - Rosswog & Bruggen - Chap 4 Evolution of stars with mass M >8Mʘ 1/4 - H-R diagram M >8Mʘ Fusion Ne, O, Si
.. Ni, Fe 4Mʘ4Mʘ the C-O core Mass > Chandrasekhar Mass
Fusion C Evolution of stars with mass M >8Mʘ 2/4 - Succession of nuclear reactions Fusing H H,He Fusing He Fe,Ni ( Longair Fig.14.6 ) Fusing C Fusing O Fusing Ne Fusing Mg Fusing Si - Each stage of fusion in the core generates a shell
‘onion’ shape - Simultaneous thermonuclear reactions in several shells
High rate of energy release
expansion outer layers
‘Supergiant’ Note: Density core ~ WD density Evolution of stars with mass M >8Mʘ 3/4 - Evolutionary stages of a star with: M=25Mʘ Fusion stage Core temperature (K) Duration Hydrogen 4x107 7x106 years Helium 2x108 7x105 years Carbon 6x108 600 years Neon 1.2x109 1 year Oxygen 1.5x109 6 months Silicon 2.7x109 1 day (Star lifetime very short compared to the evolutional time scales of a galaxy) - When an element ends its fusion in the core
contraction
T increase - Each stage of thermonuclear reactions triggers the following stage - Increase in density and temperature
successive stages reactions faster Evolution of stars with mass M >8Mʘ 4/4 - End of a massive star - Once a Fe-Ni core is produced: - No further energy can be gained by nuclear fusion (production of nuclei larger than Fe requires energy) The star has run out of fuel in the centre ! No radiation pressure to balance gravity Core electrons degeneracy pressure at Chandrasekhar limit Core collapses ‘Supernova’ Depending on the mass lost, the star ends as a ‘Neutron star’ or a ‘Black hole’ Red Supergiant Star - Betelgeuse - Mass = 20 Mʘ - Radius = 630 Rʘ - Luminosity= 63000 Lʘ - Temperature= 3500 K Hubble Space Telescope - Given the size and proximity of this Supergiant star, it has one of the largest angular diameters as observed from the Earth Supernova Explosion Type II 1/9 - When M >8Mʘ
Core-collapse supernova explosion - Theoretical models predict different phases: 1) Let’s start from the end of the massive star evolution: - ‘Onion-layered’ shells of fusing elements: - Iron and Nickel core formed: Wikipedia Supernova Explosion Type II 2/9 2) Core collapse: - No thermonuclear reactions in iron core:
no force anymore to prevent the collapse - Core mass > Chandrasekhar mass:
electrons Fermi pressure not enough
Core contraction:
Rapid heating
T=5x109K
Photodisintegration: - High energy photons split Fe into 4He nuclei and neutrons - Energy absorbed
pressure decrease
further collapse (1/4 sec, it took a few 106 years to build up core) Wikipedia Supernova Explosion Type II 3/9 3) Compression of the inner core: - Electron forced to combine with protons: e − + p → n +ν - Inverse β decay
Production of large number of neutrons and neutrinos Wikipedia - Removing electrons
pressure decrease
Gravity overwhelms electron pressure
further collapse - Densities so high that:
ν cannot escape immediately
production of a degenerate neutron gas with associated pressure The whole free fall core collapse happens in a fraction of second until a neutron star is formed in the centre (*) Supernova Explosion Type II 4/9 4) The core bounce: - If the core mass MCore < 3Mʘ (*):
the collapse stops due to the neutrons Fermi pressure - The neutron star core is stiff:
material still falling towards the centre bounce back
Core bounce:
extremely intense Shock wave propagating outward Wikipedia Supernova Explosion Type II 5/9 Wikipedia 5) The shock wave: - The collision between the shock wave and the very fast inward falling material stalls temporarily the expansion: 10-20ms (the shock wave breaks infalling iron nuclei into p and n) - How does the explosion proceed then ?
Probably additional energy is provided by the huge amount of neutrinos escaping from the core
Some neutrinos interact with stellar layers depositing energy
Re-acceleration of the shock-wave
explosion within hours or days Supernova Explosion Type II 6/9 6) Final explosion: - Several questions are still unsettled - From computer simulations, if the shock wave is: Wikipedia - Isotropic:
wave energy absorbed by outer layers - Anisotropic:
convection and turbulence
the shock wave can reach the star surface
material ejected in irregular clumps Surrounding material blasted away leaving only a degenerate remnant Supernova Explosion Type II 7/9 - Summary: - The whole process liberates ~ 1046 Joules !
~ the binding energy of the newly formed neutron-star ! - Most of the energy carried away by neutrinos that escape in a short burst (10% of the star rest mass, ~10 sec) - Only 1% energy released as e.m. radiation: - since the material is optically thick at the beginning
the radiation cannot escape immediately First sign of Supernova are the neutrinos - In the blast of the explosion all the elements produced inside the star are sent into space
the iron in our blood was produced in this way ! Supernova Explosion Type II 8/9 The sky in 26Al C-GRO satellite (Half-life ~700000 years) The decay of long-lived elements, e.g. 26Al and 56Co, is observed in the gamma-ray as nuclei emission lines - Note - The SN explosion provides energy for new thermonuclear reactions:
production of more chemical elements heavier than Fe: Au, U, ... Supernova Explosion Type II 9/9 Abs Mag - Light curves -20 Type Ia -18 Type II -16 -14 0 100 200 300 Time (days) SN type II light curves are variable because the progenitors are different each time (different peak brightness and decay times) - Typically: M v ≈ −17 Extragalactic Supernovae - SN1987A - Supernova in the ‘Large Magellanic Cloud’ dwarf galaxy - The first bright Supernova observed with modern telescopes - The brightest since Kepler’s supernova in 1604 - It was visible with the naked eye for months in the Southern Hemisphere Supernova SN1987A 1/3 - Light curve (bolometric luminosity) First 200 days - In the first 20 days: radiation primarily due to heat of the shock wave on star outer layers - Maximum after 85 days, then slow decline First 5 years Outburst Exp decay - As the expanding gases cooled: radiation from the decay of radioactive isotopes produced in the explosion (Co, Ni, Ti) - Exponential decay half-life: 77 days Supernova SN1987A 2/3 - SN1987A was a peculiar Supernova ( Longair Fig.14.11 ) - Peak luminosity was a tenth of the usual value: 108 Lʘ - The progenitor was a Blue Supergiant (Population II: young stars with low % metals) The star initially was a 20Mʘ ’Red Supergiant’ that had a strong mass loss and moved in the blue region The core collapsed when the progenitor was a 16Mʘ ’Blue Supergiant’ Supernova SN1987A 3/3 - Observations of neutrinos A 12 seconds burst of neutrinos was detected by two experiments: Kamiokande (Japan) & IMB (USA) Total of 20 neutrinos in the energy range 6-40 MeV In 12 seconds SN1987A emitted 1058 neutrinos ! Neutrinos were detected 3 hours before the explosion was observed optically Neutrinos emitted during core collapse Optical light diffused and emitted in the following processes Supernovae and their remnants - Introduction - Evolution of low mass stars - White dwarfs - Evolution of high mass stars - Supernovae Type II - Supernovae Type Ia - Supernova Remnants (SNR)