FIA0221: Taller de Astronomía II Lecture 14 Spectral Classification of Stars
Spectral types along the stellar CMD. Oh, Be A Fine Girl Kiss Me!
Classification of Stellar spectra: The MK system: strong He+ lines, no H lines O star T>30.000K strong HeI lines, weak H lines B star T~ 11.000 30.000 K strongest H lines, weak Ca+ lines A star F star T~ 7500 11.000 K T~ 5900 7500 K H grows weaker, Ca+ stronger weak metals begin to emerge strong Ca+ and Fe+ dominate, H disappearing G star T~ 5200 5900 K strong metal lines lines, weak CN CH molecular bands K star M star L star T~ 3900 5200 K T~ 2500 3900 K T~ 1300 2500 K strong TiO and VO molecular bands also neutral metal lines strong molecular absorption bands of neutral metals. But not TiO or VO anymore
The Spectral Classification of Stars
The Spectral Classification of Stars Template Spectra for Classification can be found in:.../fia0221/clases These are ascii files that can be overplotted (with SM) on top of the stars that you want to classify.
The Stellar Continuum Remember that we did not perform any Flux Calibration. Therefore the shape of the continuum is not the real shape of the stellar continuum, but the convolution between this and the spectrograph sensitivity curve. Do not use the continuum shape to classify the stars observed from Santa Martina.
Periodic Table of elements
Spectral types where the change in these lines is useful for classification More lines in: http://www.astro.puc.cl/~mzoccali/cursos/fia0221/clases/commonlines.ps
The MK system:
Luminosity classes
Luminosity classes Different luminosity classes (for the same spectral type) can be distinguished from the spectra. Even if the Teff is the same, the pressure at the star surface (hence the gravity, logg) is different. Low density (giant) stars have narrower lines, if compared with high density (dwarf) stars of the same spectral type. supergiant giant dwarf
The color magnitude diagram The HIPPARCOS CMD
The emission spectrum of the sky (earth atmosphere) This is an example of a sky spectrum. The intensity (and nature) of each emission line depends on many factors such as the presence of the Moon, local humidity and dust... + the light pollution from nearby city lights.
The absorption spectrum of the sky (earth atmosphere) wavelength (Angstrom)
Analysis of Stellar Spectra: from line strength to atomic abundances Remember how absorption lines are formed:
ATOMIC ABUNDANCES: Analysis of Absorption Lines Abundances of Atomic Elements (neutral or ionized) can be measured from the area of absorption lines. Equivalent width : width that the line would have if it were rectangular, from the continuum to zero. Strong lines may reach saturation: all the light is absorbed in the central wavelength (the bottom of the line goes to zero) and the width of the line increases very little with increasing abundance.
Analysis of Stellar Spectra: from line strength to atomic abundances We see that, in the Sun, Calcium lines are stronger than Hidrogen Lines. We know that: the hydrogen lines are produced by H atoms jumping from n=2 to n=3 the calcium lines are produced by Ca ions jumping from n=1 to n=2 The reason for the Ca lines to be stronger than the H lines can be one of the following: 1. There are more Ca+ atoms than H atoms in the solar photosphere and/or 2. Lots of Ca ions are in level n=1, whereas very few H atoms are in level n=2 and/or 3. there are lots of photons with just the right energy to excite Ca ions to n=2, but few photons with just the right energy to excite H atoms to n=3
Analysis of Stellar Spectra: from line strength to atomic abundances If the gas is collision dominated (atoms are primarily excited by collision with other particles) and if the atoms are immersed in a gas in thermal equilibrium (velocity obeys Maxwel Boltzmann statistics) then the ratio of the population of two levels of an ion is given by Boltzmann equation: Nn gn = exp Nm gm n m ( ) KT where gi = statistical weight = 2J+1 (where J = inner quantum number) gives the degeneracy of a given level (e.g. there are g=2 ways that the e of a H atom can be arranged in level 1, but g=8 ways that they can be arranged in level 2, etc... these values are tabulated ) i = excitation potential (energy between the ground level and the i th level) With this equation, given the T of the gas, we can establish how many atoms there are in each excited state, and so start finding the right answer among the previous hypothesis.
Analysis of Stellar Spectra: from line strength to atomic abundances Boltzmann equation says that the higher the temperature, the higher the excitation level of a given ion. This explain why going from spectral type K to A (i.e., increasing the surface temperature from 4,500K (K) to 10,000 K (A), the Balmer lines get stronger and stronger. However, going from spectral type A to O, Balmer lines decrease in intensity, even if 1) the temperature is still increasing 2) the H abundance is about the same why is that?
Analysis of Stellar Spectra: from line strength to atomic abundances We need to consider the effect of ionization. If all the H atoms are ionized and the T is too high to allow recombination, then there are no bound electrons than can produce Balmer lines. We then need to know, at a given T, how many H atoms are neutral (even if excited) and how many are ionized. This is given by Saha equation: N1 2 u1(t) Pe = (2 m)3/2 (KT)5/2 exp N0 h3 u0(t) where: ( I ( ) KT ) i un (T) = partition functions = gn exp KT I = excitation potential (energy between the ground level and the continuum) Pe = electron pressure
Analysis of Stellar Spectra: from line strength to atomic abundances In practice we use a model atmosphere that, applying the former equations, summarizes the characteristics of the surface of a star with a given temperature and pressure. This model tells us how to convert from measured Equivalenth Width to Atomic Abundance, by construction of the so called Curve of Growth: reduced E.W. a = damping constant 0 = column density linear region saturation region damping region
ATOMIC ABUNDANCES: Analysis of Absorption Lines So far we assumed that T and P are known from some other source (usually photometry). However, given a rough estimate of the T and P, one can iterate on spectroscopic analysis to determine simultaneously the elemental abundance and a spectroscopic T and gravity. For instance. In stars like the Sun, there are several Fe lines. A Fe abundance can be derived, independently, from each of them. EXCITATION EQUILIBRIUM If the model atmosphere is correct, then lines of FeI with different excitation potential must give the same Fe abundance. If they don't give the same abundance, something is wrong in the assumed Temperature. One can modify slightly the assumed T until all the FeI lines converge on the same abundance. IONIZATION EQUILIBRIUM On the other hand, if the model is correct, lines of Fe in different ionization degrees (FeI vs FeII) must give the same Fe abundance. Since FeII lines are very sensitive to pressure, while FeI lines are not, then the ratio FeI/FeII is used to estimate a spectroscopic pressure (gravity). NOTE: All this is true for Iron in the Sun. The dependence of each element on T and P in different stars has to be evaluated case by case.
The spectrum of hot nebulas H [OIII] 4959.5 [OII] 3727 H [NII] 6584 H [OIII]4363 HeI 5876 [SII]
The spectrum of hot nebulas H [OIII] 4959.5 [OII] 3727 H [NII] 6584 H HeI 5876 Depending on the physical conditions of the nebula, some of these lines are more sensitive to temperature and some others to pressure/density. With the help of theoretical models, one construct a diagnostic diagram, showing the run of each line with T and Ne: [SII]
Tarea 5 Entrega 29 Noviembre Introduction Very brief description of a spectrograph. Main characteristics of a spectrum. Observations Description of the observations. What was the spectrograph configuration. Describe all the exposure taken, how many frames (why so many?) and the exposure times (why this exptime?). Data reduction Describe the reduction of the data. Avoid to list all the keys entered and the directory names, etc. Explain what was done conceptually. For Example: All the dispersion lines between 5 and +5 px around the mean trace were summed (averaged?) using no rejection (or the xxxx rejection algorithm). As background, a region of [ 40: 20,+20:+40] was selected. Pixels inside this regions were summed (or averaged or medianed) with the yyy rejection algorithm in order to avoid... Data analysis Which lines were identified in each spectrum, how big they were (EW) and how was the star classified. Include also classification of the stars in http://www.astro.puc.cl/~mzoccali/cursos/fia0221/data/stars_to_classify.tar.gz Discussion and Conclusion Spectral Type of the observed stars (and of the literature ones). Any other consideration (e.g.: we should have exposed longer in star HDXXXXX because the S/N was very poor).