Lecture 19. Stellar Classification and Evolution Spectral Classification All Spectra of Stars Show Absorption Lines, Where Eleme

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Lecture 19. Stellar Classification and evolution Spectral Classification All spectra of stars show absorption lines, where elements and ions in the photospheres of those stars absorb the radiation coming from the deeper layers. Stars like the Sun show many absorptions lines due to both neutral and lightly ionised elements. Some stars show evidence of only light elements such as hydrogen and helium. Other stars have the spectral signatures of molecules in their atmospheres The first classification was due to the strength of Hydrogen lines in stars

A, B, . . . O. Strong Weak →

It was subsequently realized that the absorption lines in stars de- noted a temperature sequence.

OBAFGKM Oh Be A Fine Girl/Guy/Gerbil, Kiss Me! Class Temperature Spectralfeatures O > 30, 000K IonizedHelium,weakHydrogen

B 10, 000 30, 000 K Neutral Helium, Hydrogen lines → ionized metals (CNO) A7, 500 10, 000 K Hydrogen strongest, helium disappeared → singly ionized metals F6, 000 7, 500 K Weaker Hydrogen, neutral → & singly ionized metals G5, 000 6, 000 K Ca II most prominent, weak Hydrogen → neutral metals K3, 500 5, 000 K Neutral metals, some molecules → (CH, TiO) M < 4, 000K TiOdominates

Variations in the absorption line widths correlate with a stars luminosity and absolute magnitude MV Thus each star can be assigned a luminosity classiﬁcation. Class Name

I Supergiant II Bright giant III Giant IV Subgiant V Dwarf

The temperature classiﬁcations are further subdivided by number, with 0 being the hottest, 9 being coolest. Examples of spectral classiﬁcation: Vega A0V Aldebaran K5III Sun and Alpha Centauri G2V Sirius AIV

In the last decade, infrared surveys have discovered a class of stars cooler than M. These show complex molecules and even water in their atmospheres, and are called L and T dwarfs. Star Formation (a) Collapse The collapse of interstellar gas clouds under gravitational free-fall is believed to cause star formation.

Consider a cloud of mass M, initial radius R and density ρ0. To calculate the free-fall time, follow a test particle falling from the outer edge of the cloud towards the center, so that it follows an elliptical orbit of semi-major axis a = R/2.

4 32 M = πR3ρ = πa3ρ 3 0 3 0 . Kepler’s 3rd law states that

P 2 4π2 = a3 GM

2 3 2 12π a 3π P = 3 = 32πa ρ 8Gρ0 The free-fall time will be half of this, or

3π 1/2 tff =( ) 32Gρ0 The typical densities of interstellar clouds are 104 Hydrogen ∼ 3 5 atoms cm− , giving t 10 years. ff $ (b) Ignition As the core heats up, a point is eventually reached when nuclear fusion occurs.

1H+1 H 2D+e+ + ν 1 1 → 1 1H+2 D 3He + γ 1 1 → 2 3He +3 He 4He +1 H+1 H 2 2 → 2 1 1

The result of this reaction is that 4 hydrogen atoms create one helium atom with an associated release of energy.

The central temperature of the star depends on the mass M. An object with M 0.08M will not become hot enough for ≤ & H-fusion to begin. Once the collapse is halted by the thermal pressure of the gas, the star will gradually cool. These “failed stars” stars are called Brown Dwarfs (BDs). Real BD’s have Lithium in their atmospheres, whereas most stars do not (some very young stars show lithium). (c) Stabilization The beginning of nuclear fusion in the core of the star creates a large increase in the optical luminosity of the star that literally blows the remaining gas away from it. Such newborn stars are easily seen in nearby star-formation regions and are known as T-Tauri stars. When the radiation pressure from the core nuclear reactions bal- ance the gravitational collapse, the star is then on the main-sequence.