Supernovae: Dying Stars

Star Death

Lifetime of a Star

It is true that all living things come from stardust. In about 5 billion years, our Sun will have swelled to a red giant and engulfed the inner planets, ready to explode in a supernova. Supernovae enrich the interstellar medium with high mass elements, like iron and calcium. The high energy from supernovae also triggers formation of new stars. On average, supernovae occur only about once every 50 years in the Milky Way Galaxy. They are rare events— so rare that the last one in the Milky Way was discovered in 1604 (SN 1604, or Kepler’s Supernova)— spectacularly luminous and extremely destructive. In fact, supernovae can cause bursts of radiation more luminous than entire galaxies and emit as much energy as the Sun will in its entire lifespan! In a supernova, most of the star’s material is expelled into space at speeds up to 30,000 m/s. The shock wave passes through the supernova remnant, a huge expanding shell of gas and dust. Supernova are caused either by the sudden gravitational collapse of a supergiant star (Type I Supernova) or a white dwarf accreting enough mass or merging with a binary companion to undergo nuclear fusion (Type II Supernova). White dwarfs are very dense stars that do not have enough mass to become a neutron star (formed from supernova remnant, stars comprising almost entirely of neutrons). Supernovae can be used as standard candles (objects with known luminosity). For instance, the dimming luminosity of distant supernovae supports the theory that the expansion of the universe is accelerating. Now, with powerful telescopes like Hubble, many supernovae are discovered each year. How perfectly supernovae represent the circle of life: from death comes life!

History of Supernova Observations (Milky Way)

  • SN185 by Chinese astronomers
  • SN1006 by Chinese and Islamic astronomers
  • SN1054 (caused Crab Nebula)
  • SN1572 by Tycho Brahe in Cassiopeia
  • SN1604 by Johannes Kepler

* Supernova (SN) are named by the year they are discovered; if more than one in one year, the name is followed by a capital letter (A, B, C, etc.), and if more than 26, lowercase paired letters (aa, ab, etc.) are used

Below is a video on supernovae! Enjoy.

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Temperature of the Universe

What is the universe’s temperature? How has it changed and evolved? What causes the temperature to change? How is the temperature estimated? Is it continuously cooling or constant? –Pcelsus

Black Body Curve of the Cosmic Microwave Background

13.75 billion years ago, the Universe was much smaller and hotter. In the 1960s, Robert Dicke predicted a remnant “glow” from the Big Bang. In 1965 at the Bell Labs, radio astronomers Amo Penzias and Robert Wilson discovered that glow, named the cosmic microwave background radiation. The CBR was seen in all directions in empty space, with a black body curve (temperature ~3K in every direction). About 1 second after the Big Bang, the Universe was very hot, at ~1 billion K. At 3 minutes, protons and neutrons combine to form the nuclei of atoms. As space cooled, material condensed and atomic particles, then elements, molecules, stars, and galaxies formed. The hydrogen/ helium ratio (3:1) found today is about the same as what’s expected after the Big Bang. Atoms were “ionized” with electrons roaming free without being bound. At 300,000 years after the Big Bang, the Universe becomes transparent with a temperature of 3,000K. Light red-shifted by a factor of 1000, and the expansion of the Universe ensued.

Today, the Universe is 2.73K, or 2.73°C above absolute zero, but at the beginning of space and time, the Big Bang, the Universe reached over one billion degrees. From a single pinpoint, the Universe emerged as a scorching hot primordial soup of subatomic particles moving at high velocities. As the Universe expanded, the temperature cooled as more space was created and density decreased. The Universe is continuously cooling as it expands.

Measuring the temperature isn’t as simple as sticking a thermometer in space and waiting until it stabilizes at a certain temperature. Instead, scientists measure indirectly using the cosmic microwave background, or leftover radiation emitted by hot plasma 38,000 years after the Big Bang. As the Universe expanded, the electromagnetic waves of the CMR elongated and decreased in energy, leading to cooler temperatures. Using Planck’s law, scientists measured the black body radiation of the Universe. Planck’s law states that every object radiates electromagnetic energy according to temperature. Black body curves are lopsided, with the curve peaking at different wavelengths depending on the object. In fact, space has a nearly perfect black body curve, since physical objects tend to absorb and reflect light in certain wavelengths.

Meteroite Dating

How do scientists determine the age of meteorites, most of which are around 4.5 billion years old. –Amy

Meteorite

Meteorite Dating

Scientists measure the age of meteorites with the decay of radioactive isotopes. What is an isotope? Isotopes are elements with the same number of protons but a different number of neutrons. For instance, carbon-12 has 6 neutrons and carbon-14 has 8 neutrons. Some isotopes are very unstable and tend to decay into lighter elements by alpha or beta particle decay. Scientists use the half-life of certain elements to date objects. First, scientists must determine to isotope to use by examining the elemental composition of the object. For meteorites, scientists generally use Rubidium-87/ Strontium-87 decay, which has a half-life of 49 billion years. Rubidium-87 decays into Strontium-87. So if the object has 50% Rubidium-87 and 50% Strontium-87  (only formed by decay process), then the object is 49 million years old. Since some Strontium-87 may have been present originally, scientists use Strontium-86, whose content remains the same, as a reference. Determining the ratio between Rubidium-87/ Strontium-86 and Strontium-87/ Strontium-86 via mass spectrometer (vaporizes a tiny portion of the meteorite to form ions; the ions are then separated by mass in a magnetic field), scientists can then calculate the amount of each isotope present in the meteorite, and thus the age of meteorites. Although radioactive dating is the best method for scientists to date meteorites, many factors, such as the amount of sunshine or heavy rain, can affect measurements.

Nuclear Fusion: What Fuels Stars

CONTRACTION OF PRE-MAIN SEQUENCE STARS

  • The interior heats due to gravitational contraction and radiates away this energy as black-body radiation
  • At 10K, fusion starts, pressure increases, and the star establishes hydrostatic equilibrium (the balance between gravity and gas pressure)
  • As gravity pulls inwards (fusion releases energy, and maintains the core’s high temperature), gas pressure pushes outwards (high temperature prevents the star from collapsing under its own weight)
  • When a star reaches hydrostatic equilibrium, it enters main sequence

* Energy produced more efficiently at core’s center

Difference Between Fission and Fusion

Nuclear Fission vs. Nuclear Fusion

Fission: splitting heavy nuclei into lighter ones (e.g. atomic bombs and nuclear reactors derive their energy from fission of uranium or plutonium)

Fusion: merging light nuclei into heavy nuclei (e.g. how stars shine, hydrogen bombs, “nuclear burning” – different from ordinary chemical burning processes)

Strong Nuclear Forces: protons in the nucleus repel by electrical forces, but strong nuclear forces, which can only occur at close distances, keep the atom together. As temperature rises, protons move faster. When 2 protons fuse, the output is 1 neutron, 1 positron, and 1 neutrino.

How Fusion Works: Proton-Proton Chain & CNO Cycle

Common Elements (and Their Isotopes) Involved in Fusion: ¹H (hydrogen) [1 proton], ²H (deuterium) [1 proton, 1 neutron], ³H (tritium) [1 proton, 2 neutrons], ³He (helium-3) [2 protons, 1 neutron], 4He (helium-4) [2 protons, 2 neutrons]

Proton-Proton Chain

Proton-Proton Chain

Step 1: 2 hydrogen nuclei –> deuterium nucleus => releases positron + neutrino

  • Positron (e+): antimatter of electron
  • Neutrino (ν): unchanged particle that only interacts very weakly with normal matter

Step 2: deuterium + hydrogen nuclei –> helium-3 => releases gamma ray

-> Repeat first two steps.

Step 3: 2 helium-3 –> helium-4 => releases two protons

Summary

Input: 6 protons

Output: 2 positrons, 2 neutrinos, 2 gamma rays, 1 helium nucleus, 2 protons

Net Output: 4 protons –> 1 helium-4 => releases 2 positrons, 2 neutrinos, 2 gamma rays

0.7% of the total mass of 4 protons is converted into energy, while 99.3% results in 1 helium nucleus. Some of the mass is converted into energy. Since E = mc², a little mass and release tremendous energy. While at rest, however, energy is equal to mass.

CNO Cycle

CNO (Carbon-Nitrogen-Oxygen) Cycle

The CNO Cycle is the main nuclear burning chain in main sequence stars hotter than the Sun. Using carbon as a catalyst to convert hydrogen into helium, the CNO cycle also converts 7% of hydrogen’s mass into energy; hydrogen fuses with carbon to form helium. 10% of the Sun’s nuclear fusion reactions is from the CNO Cycle. In 1967, Hans Bethe theorized on the energy production in stars.

Interstellar Medium – The Material Between Stars

WHAT LIES BETWEEN STARS IN GALAXIES?

– Interstellar Medium

Interstellar Medium

Interstellar Medium is gas and dust between stars, nebulae, and giant molecular clouds (basic building blocks of galaxies in star formation). The four types of matter in interstellar medium are: interstellar dust, interstellar atoms, interstellar molecules, and interstellar snowballs.

Interstellar Dust

  • Interstellar Reddening: dust that scatters blue light and causes stars to look redder
  • Extinction of Obscuration: high dust content that diminishes the brightness of stars, by as much as 25 magnitudes
  • Can be smaller than smoke particles
  • Consists of graphite, silicates, or ices
  • In core of heavy elements (e.g. iron, magnesium), mantle of organic compounds (oxygen, carbon, nitrogen), and outer mantle of ice

RADIO ASTRONOMY

  • Radio waves = longest wavelength of electromagnetic waves
  • Brightest optical objects not necessarily the brightest radio objects
  • e.g. Taurus A (Crab Nebula) and Sagittarius A (center of the Milky Way Galaxy)
  • Radio Spectral Line: the frequency or wavelength at which radio noise is slightly more or less intense
    • Hydrogen: 21 centimeter line
    • Radio spectra lines of molecules
      • OH (hydroxide): 1963
      • H20 (water): 1968
      • NH3 (ammonia): 1968
    • Over 50 molecules in interstellar space
    • Gives information on temperature, density, and motion
    • Molecular absorption line in UV

Interstellar Molecules

  • Molecules: two or more atoms bound together (e.g. H2O, CO, CH4, OH, H2, NH3)
  • Give absorption or emission bands
  • Observable in very cold, low density interstellar environments

Interstellar Snowballs

  • Between the sizes of  grains and comets
  • Composed of water, carbon, silicates, and other molecules

Interstellar Regions

  1. HI region: 200 K
  2. HII region: 10,000 K
  3. Molecular clouds: 50% gas in our galaxy
  4. Hot interstellar medium: 1 million K, super-heated gas from expanding supernova blasts (up to 90% of total volume)
  • HI Region
    • High density of neutral hydrogen atoms about a million atoms per cubic centimeter (e.g. Orion Nebula)
    • ~ 200 K
  • HII Region
    • Hydrogen with electron removed; e.g. ionized hydrogen gas (in emission nebulae)
      • Average density of hydrogen elsewhere is 1 atom per cubic centimeter
    • ~ 10,000 K

What are White Dwarfs?

White Dwarf

White dwarfs are the bare cores of low-mass stars such as the Sun. A low-mass Main Sequence star becomes a white dwarf when the star uses up all its hydrogen, swells, and ejects its outer layers.

If a white dwarf has a binary companion…

  • Mass-Transferring Binary Star System (e.g. white dwarf and red giant)
  • Gas “spills over” from red giant’s atmosphere and is gravitationally pulled into the white dwarf
  • Nova– explosion powered by fusion of hydrogen to helium on the surface of a white dwarf star; caused by matter spilling onto the star from its binary companion
    • Star brightens rapidly, then fades over weeks or months
    • Nova explosions can recur in the same binary system

Maximum Mass

  • S. Chandrasekhar(1930): calculated the maximum mass of white dwarf
    • Electron degeneracy pressure can only support a white dwarf less than 1.4 M☉
    • If a white dwarf accreted enough mass that overcomes this limit, gravity would win and something dramatic would happen
    • Chandrasekhar’s Limit: a white dwarf’s mass can only be less than 1.4M☉

After the Type Ia Supernova, the white dwarf is completely destroyed; no solid remnant is left, although the companion star might remain.

Type Ia Supernovae

  • Brightens over 2 weeks, reaches a peak, and then fades
  • At its peak, the supernova is 10 times more luminous than the Sun (e.g. 1994 D in galaxy NGC 4526)
  • Composed of mostly iron and other heavy metals
  • Core collapse supernova produces carbon, oxygen, neon, magnesium, silicon, and other lighter elements, and iron and other heavy elements
  • The ejected material is “recycled” into new generations of stars and planets

*Note: Without supernova explosions, earth-like planets, organic chemistry, and life wouldn’t exist.

  • All heavy elements were created inside stars or during supernova explosions, and then expelled into interstellar space
  • Heavier elements in supernova form by neutron capture: in an dense environment of  free neutrons, atoms absorb neutrons, beta (β) decays, and a proton forms — when an atom gains a proton, its identity changes and it moves one atomic number on the periodic table
  • Carbon, nitrogen, and oxygen are winners in the burning (release tons of energy)
  • Lithium, beryllium, boron are destroyed and not created in stars
  • Iron, the most stable element, is the end of nuclear burning
  • More massive elements formed by neutron capture followed by β-decay