Download SUPERNOVA EMISSION FROM GAMMA RAYS TO THE INFRARED

SUPERNOVA EMISSION FROM GAMMA RAYS TO THE INFRARED
Anders Jerkstrand, Stockholm University
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Supernovae
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A supernova is the explosion of a star, which occurs about once per century per galaxy.
When they occur, they often outshine the entire galaxy for a few weeks.
SN1994D in NGC 4526
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Supernovae
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These stellar explosions can be of two types
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Gravitational collapse of a massive star (M > 8 Msun).
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Neutron star or black hole forms.
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Labelled Type II, Ib or Ic.
Thermonuclear explosion of a white dwarf approaching the Chandrasekhar mass.
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No remnant.
Labelled Type 1A.
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Supernovae
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In both cases , about 1051 ergs is liberated in the form of radiation (~1%) and kinetic energy (~99%). This gives expansion velocities of ~104 km/s, giving broad spectral lines.
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Supernovae
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This energy (1051 ergs) agrees well with the thermonuclear energy of burning ~ 1.4 Msun of carbon/oxygen to iron. (~0.001*Mc²)
But the binding energy of a neutron star is >~1053 ergs. (3/5GM2/R with M=1.4 Msun and R=10 km). Where does that energy go?
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Supernovae
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In 1987, Kamiokande II, IMB and Baksan gave us the definite answer by registering 24 neutrinos from SN1987A, three hours before the explosion was observed in electromagnetic bands.
Some 99% of the 1053 ergs released in the neutron star formation emerges as neutrinos.
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Kamiokande II
Outbreak
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As the shock wave breaks through the surface, there is an initial and brief (~minutes) burst of hard radiation at UV and X­rays.
The only such outburst observed (!), lasting 6 minutes, occured in 2008 as Swift was pointed in the right direction at the right time (SN2008D).
X­rays
Optical a few days
later
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Lightcurve
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After the breakout, the supernova cools quickly and behaves as an expanding fireball at ~5000­10,000 K.
The luminosity increases as the surface area of the photosphere grows, then peaks and falls off as the gas becomes optically thin.
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Lightcurve
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After a few weeks or months, the light­curve is seen to settle down into a strict exponential decay L ~ exp(­
t/) with  ~111 days.
Since this matches the decay time of the radioactive 56Co, Colgate & McKee suggested in 1969 that this isotope powers supernovae in this phase.
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Nucleosynthesis
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Modern hydrodynamical explosion models confirm that substantial amounts of radioactive elements are created in the explosion, the most important being 56
Ni (which quickly decays to 56Co), 57Ni and 44Ti.
Explosion model of
A 25 Msun star from
Woosley & Weaver 1995.
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THE ENERGY INPUT BY THE THREE DOMINANT RADIOACTIVE ELEMENTS
Power [ergs/second]
Total power Using the masses determined
From SN1987A:
M(Co56) = 0.07 Msun
M(Co57) = 3.3*10­3 Msun
M(Ti44) = 1*10­4 Msun
Co56
Co57
Ti44
Co56­phase
Co57­phase
Ti44­phase
Time [years]
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Radioactive decays
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The decays usually occur by electron capture, ­ (electron) or + (positron) channels to an excited state in the (A,Z+1), (A,Z­1) daughter nucleus.
This state deexcites by the emission of gamma rays, with energies often exceeding the first step.
Ex for 56Co:
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3% of the energy emitted in positrons.
97% emitted as gamma rays (mainly at 847 and 1238 keV). 12
Gamma rays
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The gamma rays Compton scatter off free and bound electrons.
From the Klein­Nishina cross section it can be seen that the ejecta becomes optically thin to the gamma rays after a few years.
Detection of the gamma
rays
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Hard X­rays, being Compton down­scattered gamma rays from 56Co decay, were detected from 1987A by the Ginga satellite and by the Kvant module on Mir already after ~6 months.
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Detection of the gamma
rays
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Soon, also direct detection of the 56Co gamma rays (at 847 and 1,238 keV) from SN1987A were made by the Solar Maximum Mission (SMM).
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Detection of the gamma
rays
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The detection of X­rays and gamma rays occured much too early for the 56
Co to reside deep in the core, protected by the envelope.
The conclusion was that mixing has occured where 56
Co­clumps have been propelled far out into the envelope.
Carbon
Oxygen
Nickel
3D explosion of a 16 ZAMS progenitor.
Three components displayed.
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Detection of the gamma
rays
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For the 44Ti gamma rays, the emission from 1987A has been too weak to be detected.
However, in 1994 COMPTEL detected the 1157 keV line from 320­year­old supernova remnant Cas A. Later, the 68 and 78 keV lines were detected with BeppoSAX.
.
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Supernova model spectra
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In Stockholm, part of the work is aimed at computing supernova model spectra in late phases (t >~ 1 year), and comparing those with observations.
To do that, a large set of coupled, non­linear equations have to be solved.
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Energy degradation
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The first step is to compute how the gamma rays (and the fast electrons and positrons) produce a set of fast ”secondary” electrons that in turn cause heating, excitation and ionization in the supernova gas.
Fast electron hits free electron
→ HEATING
Fast electron hits internal transition in
atom → EXCITATION
+
Fast electron hits collisional ionization
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channel in atom → IONIZATION
Thermal and statistical
equlibrium
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Next, the temperature, the ionization balance and the non­LTE level populations are solved for in each zone by setting up the equations for thermal and statistical equilibrium.
These equations also depend on the internal radiation field, which is solved for by doing radiative transfer with a Monte­Carlo technique.
For a 1D explosion model, typical run times to convergence are ~ 1 day on a 100 CPU parallell computer.
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Radiative transfer
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The supernova gas is in a Hubble flow just like the galaxies in the universe.
A photon is continously redshifting with respect to the co­moving frame.
5000 Å
5100 Å
EXPLOSION CENTER
Gas parcels
5150 Å
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Radiative transfer
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104 km/s
The fact that each photon has to pass many lines (all that it redshifts past before reaching the edge of the nebula) makes line absorptions very important in supernovae.
Ex : a photon emitted at 3000 Å redshifts to 3100 before reaching the edge of the nebula at ~ 104 km/s. The line opacity is the sum of all lines between 3000 and 3100 Å.
3000 Å
3100 Å
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Radiative transfer
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Especially in the UV, the iron­group elements have thousands of lines that together form a formidable opacity.
Photons at UV wavelengths are therefore easily absorbed and fluoresce to longer wavelengths.
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Radiative transfer
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This line opacity commonly makes supernovae very weak in the UV band, while at the same time adding fluorescence to the optical/NIR bands.
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The
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44
Ti-phase
After ~5 years, supernovae enter the 44Ti­
dominated phase.
The ejecta is by now largely optically thin to the gamma rays, but the positrons are efficiently trapped.
Ti has a long life­time (~86 years) and the output therefore changes slowly in this phase.
44
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The
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44
Ti-phase
The nebula is now cold (~100 K), mostly neutral, and at densities of ~ 105 particles cm­3 (t/10 years)­3.
The cooling is done by mid­IR lines such as [Fe II] 26 um. The dust is probably still optically thick and emits a blackbody spectrum of similar temperature.
All in all, about 50­70% of the deposited energy probably emerges in the mid­IR (3­50 um) in this phase.
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The
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44
Ti-phase
However, excitations and ionizations still produce emission in the UV/optical/NIR.
We have shown by modeling of the 1987A spectrum at 8 years (Jerkstrand et al, in prep.) and 20 years (Kjaer et al 2010, in press), that ~1*10­4 Msun of 44Ti is powering the deepest core regions in the supernova, producing emission in newly synthesized iron, silicon and oxygen.
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The
44
Ti-phase
Kjaer et al, 2010, in press
Observed spectrum
Model
●
Silicon
Calcium
Iron
Helium
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Dust in supernovae
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Conditions favorable for dust formation include high densities (>~ 107 cm­3) and low temperatures (<~ 1500 K). In supernovae, the temperature is initially too high, and eventually the density becomes too low.
But there is an intermediate epoch (~1­2 years after explosion) when conditions are appropriate.
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Dust in supernovae
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In 1987A, a systematic shift of the line peaks towards the blue coincided with a rapid luminosity decrease starting at day ~530.
The interpretation was that dust was forming, preferentially blocking out the far (receding) side of the nebula.
Line
profile
Frequency
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Soon, a 300 K blackbody continuum was rising in the mid­
IR, as detected by the Kuiper Airborn Observatory.
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Dust in supernovae
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Dust formation has also been observed in other (gravitational­collapse) supernovae (for example 1999em, 2004et, Cas A), usually with Spitzer.
Cas A in optical and at 24 um.
The dust temperature is 80 K.
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It appears that the amount of dust formed in supernovae is quite small (~10­4 ­ 10­2 Msun), and supernovae are currently not believed to be a major 32
production site for the dust in the universe.
Dust in supernovae
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However, it is possible that most of the dust is very cold in supernova remnants, and escapes detection by Spitzer.
Sub­mm detection by SCUBA of Cas A and Kepler have been controversially interpreted as the presence of large (~1 Msun) dust masses.
Herschel imaging of Cas A has recently verified a cold (35 K) dust component with M ~ 0.075 Msun. But modeling is tricky!
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Barlow et al, 2010 : Cas A far­IR with Herschel.
Many components, also from the ISM!
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Summary
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Supernovae are stellar explosions that produce emission in all electromagnetic bands for many decades or centuries.
Except for the first few weeks, they are powered by radioactive decays from 56Co, 57Co and 44Ti.
The computation of model spectra requires computing the degradation of gamma rays and positrons, solving the equations of thermal and statistical equilibrium, and performing Hubble­flow 35
radiative transfer.
Summary
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At late time, the supernova emits mainly gamma rays from the 44Ti decay and mid­IR emission from gas and dust at ~100 K, but also some UV/optical/NIR emission from fast­electron excitations and ionizations.
THANK YOU
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