Colophon
Stellar Radiation | |
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| In this EBook we discuss the importance of Electro-Magnetic Radiation for Astronomy | |
| Tags | Astronomy, Electro-Magnetic Radiation, spectrum, visible light, Infrared, micro wave, ultra-violet, radio, x-ray, gamma ray, Wave-Particle duality, refraction, diffraction, reflection, James Maxwell, Albert Einstein, Max Planck, photon, quantum, wavelength, frequency, speed of light, nuclear fusion, spectroscopy, spectrograph, black body radiation, star colour, absorption lines, emission lines, atom, nucleus, electron, chemical elements, stellar spectrum, line spectrum, absorption spectrum, emission spectrum, ionisation, spectral classification, luminosity, brightness, magnitude, stellar temperature, stellar mass, Annie Cannon, Ejnar Hertzsprung, Henry Russell, HR-diagram, main sequence, red giants, super red giants, white dwarfs, supernova. |
| Prerequisites | |
| Author(s) | EV |
| First Published | April 2008 |
| This Edition - 2.1 | October 2018 |
| Copyright |
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Introduction
When we look at the night sky with the unaided eye, or even through optical telescopes, we see bright stars and faint stars, and in some cases we see distinct differences in colour. But what our eyes see, is only a tiny fraction of what is called EM-radiation.
Star light, Star bright. Modern interpretation of the zodiac by Bechnokid. Modern Astronomy now has the technology to use the entire spectrum of this radiation and that has revolutionised our knowledge about the Universe.
In this module we will discuss some aspects of EM-radiation that are of importance to Astronomy, and explain how these aspects are telling us most of what we know about stars and other luminous objects in the Universe.
But first some basics...
What is Light?
Credit: recoverybridge.org Generally when astronomers talk about “light” they mean visible light to which our human eyes are sensitive and that produces a visual sensation. A more general term is Electro-Magnetic radiation (or EM-radiation) that refers not only to visible light, but to everything from Radio transmission up to Gamma radiation. One problem with a definition of this “EM-radiation” is that no scientist can give one model or theory that fully describes it. 
Credit: electromagnetic-waves.com
Since the early 19th century James Maxwell developed an electro-magnetic field theory that considers light as a wave phenomenon, one that we can imagine when we look at waves and ripples on a water surface or in an oscillating string.
Credit: guidetothecosmos.com
But Einstein demonstrated about one hundred years later that light also behaves in ways that cannot be described by a wave theory. Together with his contemporary Max Planck, he formulated the concept of the photon as a fundamental quantum of energy. Often this is confusingly referred to as a “particle” model, but photons are not “mini-jelly beans” that travel through space, but rather very small amounts (quanta) of energy.
These two interpretations of light, the wave model and the quantum model, are mutually exclusive . This is referred to as Wave-Particle Duality. One can use only one model at any time, depending on which characteristics of light one is studying. Probably the safest general term one could use for this phenomenon is “Energy radiation”, but very generally the term EM-radiation is used, also in astronomy.
Aspects of the Wave model
Credit: upscale.utoronto.ca (edited) For travelling waves we can talk about wavelength, the distance between two consecutive crests or troughs of the wave. Hence wavelength is a length and is expressed in metres. We can also count how many complete waves pass at a particular point each second and we call that frequency, expressed in cycles per second or Hertz (Hz). If we multiply the two, then we calculate what length of waves passes in one second, thus we obtain the speed of the wave in metres per second.
Light or EM-radiation travels very fast and at a constant speed (in vacuum) of 299,792,458 metres per second. This number is exact, because the definition of the metre as a unit of length has been adapted to this value for the speed of light. To have an idea how fast this is: light travels about 7.5 times around the world at the equator in one second.
With a constant speed of the travelling wave, frequency and wavelength are inversely proportional. This means that when the frequency e.g. is doubled, then the wavelength will be halved and the other way around.
Aspects of the Quantum model
If we think about the quantum model for light we can talk about the (photon) energy of light that is proportional to the frequency. In other words if the frequency is doubled then the energy is doubled as well, when the frequency is halved then the energy is halved, etc.
| In summary, when the wavelength is long, the frequency and energy are small, when the wavelength is short, the frequency and energy are large. |
The Spectrum
If we go back to visible light, we see that different wavelengths (or frequencies or energies) are identified as colour of light, and we use the term spectrum to indicate the range of the different colours from red to violet.
The table below shows the visible spectrum with its various colours and values.
| Colour |  |
 |
 |
 |
 |
 |
| Wavelength | 380–450 nm | 450–495 nm | 495–570 nm | 570–590 nm | 590–620 nm | 620–750 nm |
| Frequency | 789-668 THz | 668-606 THz | 606-526 THz | 526-508 THz | 508-484 THz | 484-400 THz |
The complete spectrum of EM-Radiation
Credit: Wikipedia
This visible spectrum is only a very small fraction of what we know as the entire EM-spectrum, from Radio waves with a wavelength of up to hundreds or thousands of kilometres, to Gamma waves with a wavelength of typically 0.01 nm or less.
All this radiation behaves fundamentally in the same way and travels through the vacuum of space at the constant speed of light.
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|
Gamma-ray
|
X-ray
|
Ultra Violet
|
Visible
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InfraRed
|
Microwave
|
Radio
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|
Energy (eV)
|
106-105
|
105-102
|
102-3
|
3.0-1.7
|
1.7-10-3
|
10-3-10-6
|
10-6-10-14
|
|
Frequency (Hz)
|
3x1020-3x1019
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3x1019-3x1016
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3X1016-8x1014
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8x1014-4x1014
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4x1014-3x1011
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3x1011-3x108
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3x108-1
|
|
Wavelength (m)
|
10-12-10-11
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10-11-10-8
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10-8-4x10-7
|
4x10-7-7x10-7
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7x10-7-10-3
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10-3-100
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100-108
|
The boundaries are not always sharply defined and are rounded off here.
Theoretically the spectrum extends until infinity on both sides.
Units used are:
- eV (Electron Volt) for Energy
- Hz (Herz) for Frequency
- m (metre) for Wavelength
For the notation of Powers of Ten and prefixes used, see the document Prefixes and Symbols in our Physics Study Notes
Only within the last hundred years astronomers can use this entire EM-spectrum of radiation that is received from celestial objects, because we now have the technology to detect and measure these different parts of the spectrum. Before that, only the visible part of the spectrum could be used.
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Often you will see that in Astronomy the unit for wavelength is the non-SI unit Ångström, named after the Swedish scientist Anders Jonas Ångström . Ten Ångström is one nm (nanometre) when we use the SI unit for length and the unit prefixes for powers of ten. The Ångström is generally not pronounced properly because the Swedish Å is pronounced as “O” not as “A”. |
Effect of the Earth’s Atmosphere
The Earth atmosphere is not transparent for all parts of the EM-spectrum. That is fortunate because the high energy part of the spectrum is lethal for bio-organisms and without the filtering effect of the atmosphere, life on Earth would not be possible.
Credit: Wikimedia
The graph indicates the transparency of the atmosphere across the spectrum. As can be seen, only the visible spectrum, parts of infrared and parts of the radio wavelengths reach the Earths’ surface and can be observed from ground-based observatories. All other parts of the spectrum can in principle only be observed from space.
The Thin Blue Line; Credit: ISS, NASA
| Thickness of atmosphere The atmosphere is really very thin. Imagine that you have a globe of the Earth with a diameter of 30 cm. The effective atmosphere of 480 km is then less than half a millimetre at that scale. Yet it protects all life on Earth against high-energy radiation from the Sun and outer space. |
Continuous spectrum
Stars generate their own energy by nuclear fusion. Based on the energy produced by that process, EM radiation over the entire spectrum is produced in the outer layers of a star through a multitude of processes and interactions.
Generally, the light we receive from a star, only tells us about what is happening in the outer layers of the star, not what happens in the inner part where the fusion process generates the energy. The visible part of the spectrum of e.g. our Sun, originates from the breakdown of a peculiar particle called the negative Hydrogen atom (ion) in the photosphere, emitting photons in the entire visible range of the spectrum, as well as in infrared.
We only need to look at a rainbow to appreciate that sunlight consists of the various colours of the visible spectrum. With a prism or a special instrument called a spectrograph (see below), we can stretch light spatially into different wavelengths, to study how the light we receive, is composed of different wavelengths or frequencies.
So generally, stars emit a continuous spectrum, although there can be differences in the intensity between different wavelengths.
| The rainbow is a natural spectrum caused by the refraction and reflection of sunlight through raindrops, and also shows that sunlight consists of a continuous envelope of colours from red to violet. Sometimes you can see a fainter outer bow caused by double reflections inside the raindrops. |
Star colour
Black body temperature curves. Credit: Wikipedia The reason that we see some stars as predominantly blue-ish in colour, and that others are more red, is because the intensity of the emitted light can be different for different wavelengths.
Ideally, a star emits light as a theoretical object called a black body, for which the intensity over the spectrum depends on the temperature of the star’s outer layer. This makes it possible to determine the temperature of that outer layer by observing the intensity of the various wavelengths in the spectrum.
A star’s blackbody curve reveals its temperature (see graph). Notice that the peak of each temperature curve, moves to the left (shorter wavelength) when the temperature goes up. This follows Wien's Displacement Law.
Astronomers use sets of standard colour filters to determine through which filter the maximum energy (light intensity) is measured. That filter (wavelength) tells us the temperature of the star’s outer layer that is called the photosphere.
Think about a red hot poker you pull out of a fire. The colour you see (which is the peak colour) reveals its temperature.
The hotter it gets the shorter wavelength light it emits. Initially you may not see any light coming off the poker, then it radiates mainly in infrared. When it gets hotter it will become red and later even white (yellowish) hot. If you use a material that does not melt at that point, you will see it become blue hot, etc.
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Famous for its display of a great variety of colours, this open cluster of stars deserves its name: The Jewel Box. One of the bright central stars is a red Super Giant, in contrast to the many blue stars that surround it. The cluster contains just over 100 stars, and is about 10 million years old. It is at a distance of about 7500 light-years and is about 20 light-years across. It can be seen with binoculars towards the southern constellation of Crux. |
Absorption lines
On closer inspection of the continuous spectrum received from stars, astronomers see that there are dark lines in it, absorption lines.
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|
Credit: woodahl.physics.iupui.edu
Atoms that contain electrons around the nucleus can have different energy states.
These energies are discrete in the sense that only very specific energy levels are possible. Atoms can change their energy state by either absorbing energy from incoming photons, or can emit energy in the form of photons. But both absorption and emission of energy only happens over discrete energy quanta.
Roll over the mouse to see the difference between the ground state and an excited state. Based on cartoon by Khalil Bendib.
Compare this with a henhouse. When everything is quiet, the hens (electrons) are in their ground state. But if they are disturbed by an outside influence (photon radiation) they can jump up to any of the energy levels where they can be. They cannot be just about anywhere, they have to be on one of the prescribed energy levels.
Credit: uwgb.edu
If light from the star shines through the outer layer, then atoms in that outer layer can absorb photon energy. This means that in the light that comes through to us, specific photon energies (wavelengths or colours) are missing or at least much less intense than other wavelengths.
Each atom has its own specific possibilities of change between energy levels and therefore each chemical element has its own specific “barcode” of dark lines, called absorption lines.
 Looking at a star we see the absorption spectrum that originates from the outer layer, against the background of the continuous spectrum. |
Emission lines
Credit: abyss.uoregon.edu
Stars generally only display absorption line spectra, but when a cloud of inter-stellar gas is excited by incoming radiation, electrons in that gas can be excited and go up in energy level through absorption. Subsequently, these electrons can come down again and thus emit photons of very specific energy. This generates bright lines in the spectrum for those specific wavelengths, an emission spectrum.
Absorption and emission lines for each chemical element are at the same location in the spectrum. The lines are either dark in an absorption spectrum and bright in an emission spectrum.
Hydrogen and other elements
Credit: solarobserving.com
The most abundant element in the Universe is Hydrogen that has only one electron in each atom. The different energy jumps of that electron however give four absorption lines in the visible part of the spectrum. Other absorption lines are in the IR and in the UV part of the spectrum.
Other chemical elements, that have more electrons form more complicated spectra, but each element produces its own specific line spectrum. For this reason astronomers can find out which elements are present in stars, even when these stars are at very great distance.
Some examples of the emission spectrum of other elements
| Hydrogen |  |
| Helium |  |
| Lithium |  |
| Carbon |  |
| Nitrogen |  |
| Oxygen |  |
| Silicon |  |
| Iron |  |
Spectroscopy
Credit: scipp.ucsc.edu (edited)
We saw that white light can be separated into different colours (wavelengths) by shining it through a prism. This uses the principle of refraction of light, that bends at the boundary between two different densities. The refraction occurs twice when light enters and exits the prism.
A better separation of the wavelengths can be obtained with diffraction, which is the bending of light when it touches a sharp edge. Diffraction is also a typical wave phenomenon and can be demonstrated with water waves and even sound waves.
| Read more here on the wave properties of light Reflection, Refraction and Diffraction. |
Modern spectrographs use a diffraction grating on which many very fine parallel grooves are etched. Extreme accuracy is required in the constant thickness of the lines and in their spacing that should be in the order of the wavelength of light.
Only very few companies can make these gratings to scientific standards, and even then the manufacturing process is one of trial and error, where only a few of the end products ever make it through the quality control. Especially for space born spectrographs, which cannot be serviced after launch, the requirements are at the limit of technical capabilities. But the results of such spectrographs are spectacular.
Schematic of a spectrograph. Credit: outreach.atnf.csiro.au (edited)
The rest of a spectrograph is fairly straightforward in principle. If only one wavelength is focused onto the detector, the instrument works as a very narrow wavelength filter (aka "monochromator"). Otherwise many but separate wavelengths are recorded on the detector. Often many individual detectors that are placed side by side each receive their own wavelength.
In early models the detector was a photographic plate, but nowadays, a CCD type camera is used, that is directly connected to a computer readout and the process is fully automatic.
Early line spectrum of the Sun recorded on photographic plate.
Wavelength-intensity graph, Keck Observatory
In observatories, the spectrograph is usually placed in a climate controlled room, away from the telescope. Often the light from the telescope enters the spectrograph through fibre optics.
No astronomer looks at the light, but only analyses the resulting spectrum at the computer, where processing can be done automatically as well.
Typical plots that are studied are not only the nice colour spectra as we know them, but more often graphs of intensity (vertical) versus wavelength (horizontal). Dips in the graph show the location of absorption lines.
A break during the installation work. Note the many individual components of the spectrograph that must be carefully aligned. At the end of this phase, they will disappear under a special cover, for dust protection and temperature stabilization. Credit: ESO
An example of a modern echelle spectrograph is the UVES instrument of the Very Large Telescope (VLT) of the European Southern Observatory (ESO) in Paranal, Chile
Signature of a star
Spectrum of Aldebaran (Alpha Tauri)
Betelgeuse. Image based on spectral image from www.weasner.com
By comparing line spectra with the spectra generated in the laboratory, astronomers can make out which elements are present in the outer layers of the star or in the inter-stellar gas. It is as if these chemical elements produce their particular “barcode” in the star’s spectrum, identifying them to the observer.
The absorption spectrum of a star shows the sum of different “barcodes” of different elements. It does not matter how far away the star is, as long as we can receive the starlight and separate it through a spectrograph, we can get this information.
If a certain element is very abundant, the absorption that takes place will be more intense. Therefore this abundance can be related to the strength of various absorption lines.
To produce certain spectral lines, a star must be hot enough to excite the electrons out of a particular state but not too hot to ionise a significant fraction of the atoms. In ionised atoms, the electrons are removed, so there are no electrons to excite, and therefore no spectral lines. Thus astronomers can also relate a star’s temperature to the strength of its various spectral lines.
High Resolution Sun spectrum
Credit: N A Sharp, NOAO/NSO/Kitt Peak FTS/AURA/NSF
This image was created from data observed with the Spectrometer at the McMath-Pierce Solar Facility at Kitt Peak National Observatory, near Tucson, Arizona, USA.
It must be read like a book: from left to right,
top to bottom.
Each of the 50 slices covers 6 nm, for a complete spectrum across the visual range from 400 to 700 nm.
Click on the image to download a high resolution version of this image (2.8 MB).
Spectral classification
Standard spectra for each class. Credit: odin.physastro.mnsu.edu
Astronomers have classified stars according to their spectra into spectral classes (also called spectral type).
The major spectral classes are type O, B, A, F, G, K, M where the O-class is the most luminous (and hottest) and the M-class is the faintest (and coolest). Each type is subdivided into 10 finer divisions (0-9), as A8 or F0.
The picture above shows standard spectra for each class.
The historical reason why this sequence seems arbitrary, is that initially the classes were defined alphabetically according to strength of the spectral lines. Later it was discovered that this is not directly related to temperature. So the original sequence was re-arranged according to temperature.
Classifying stars used to be something of an art form that comes with practice. Each spectral class is defined by the spectrum of a standard star (above), against which the other stars are compared. The classifier eventually memorizes the standards and can classify the spectrum of a random star very quickly.
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Astronomer Annie Cannon classified well over 300,000 stars in her lifetime. Annie Jump Cannon at work. |
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Modern classifiers are now using automated systems with computers and complex software, that will give more objective results.
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| Collection of photographic spectra from 1968. Credit: Wellesley Women in Science, www.wellesley.edu |
Modern optical spectrum with a zoom-in of the H-alpha line below it. Credit: http://skyserver.sdss.org |
Hertzsprung – Russell
In about 1910, two astronomers, Ejnar Hertzsprung in Denmark and Henry Russell in the US independently plotted a new type of graph, known as the Hertzsprung-Russell diagram, based on the stellar classification results at the time. Horizontally they plotted star temperature or spectral type and vertically the luminosity or energy output of the star.
In these plots they noted that stars do not appear just anywhere, but only in specific regions of the graph. Most stars appear in a band diagonally from bottom right to top left, while a minority appeared in regions above and below this band.
This Hertzsprung – Russell diagram (or HR-diagram) has become a very important instrument in modern astronomy. Since more and better information became available about distance and therefore luminosity, the diagram has become much more precise.
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It is not easy to find luminosity of a star when you don’t know its distance. We can easily measure the apparent brightness as the star appears in the sky, but that does not tell us how bright the star really is. A relatively faint star can be very luminous but at a great distance and a relatively bright star can be not very luminous but nearby. Initially astronomers used stars of one star cluster, so that there cannot be a great difference in distance. Then the apparent brightness is a measure for luminosity. They also knew distance to some nearby stars that was determined with the parallax method. |
Luminosity of an object is the amount of energy it emits per unit of time. We express the luminosity of a light bulb in Watt, which is Joule (energy) per second. For stars luminosity is commonly expressed in luminosity of the Sun (Lsun). There is a fixed relation between luminosity and absolute magnitude of a star.
Stars in the diagonal band are called main sequence stars. Appearing in this more or less linear structure in the graph means that there is a correlation between temperature and luminosity; these two properties are related. The higher the temperature of the star (further left) the larger the luminosity or power output (further to the top). This holds for stars that are in the process of fusing hydrogen to helium. The hotter stars are more massive than the cooler ones, they burn faster. So the position of any star in the main sequence gives an indication of the mass of that star.
We also see stars in the upper right (cool and luminous) and in the bottom left region (hot and faint).
But if stars of a particular spectral type (i.e. being on any vertical line in the diagram) have a larger luminosity (total energy output) this means that, according to Stefan-Boltzmann’s Law, they have a larger surface area and thus are much bigger. They are in the class of the Giant and Super Giant stars (top right above the main sequence).
Below the main sequence (bottom left) are stars that are hot but not very luminous. These are much smaller and in the Dwarf category. White dwarfs are the end stage of evolution of stars like our Sun, and are very hot and very compact.
These groupings are very important to understand the life cycle of stars and we will discuss this further in our EBook “Stellar Evolution”.
But let us now look at what the classification in the HR-diagram means for our understanding of the physical properties of stars.
The interpretation of the spectral class
Analyses of the spectra show that all the stars of the Main Sequence, those fusing hydrogen, have similar chemical compositions, all about 90% hydrogen, 10% helium, and only a fraction of other elements The differences in stellar spectra, at least for main sequence stars, are caused almost entirely by differences in ionisation (when electrons are removed from the atoms, thus producing no absorption lines) and thus changes with temperature from cool class M stars up to hot class O stars.
The image left is a modern version of the HR-diagram.
Note that the vertical scale shows both luminosity (energy output) as compared to the Sun, as well as absolute magnitude, which is luminosity expressed in the magnitude scale. (Read more about this in our EBook "Stellar Distance").
The horizontal scale at the top shows both spectral class and temperature, and at the bottom the Blue-Violet colour index.
Find a larger size of the image here.
Main properties spectral classes
The mass, radius, and luminosity listed in the table below for each class are appropriate only for stars on the main sequence. The spectral classes O through M are subdivided by Arabic numerals (0–9). For example, A0 denotes the hottest stars in the A class and A9 denotes the coolest ones. The Sun is classified as G2. (Table credit: Wikipedia).
|
Spectral Class |
Surface |
Mass |
Radius |
Luminosity |
Hydrogen |
Fraction of all |
|
O |
≥ 30,000 K |
≥ 16 |
≥ 6.6 |
≥ 30,000 |
Weak |
~0.00003% |
|
B |
10,000–30,000 K |
2.1–16 |
1.8–6.6 |
25–30,000 |
Medium |
0.13% |
|
A |
7,500–10,000 K |
1.4–2.1 |
1.4–1.8 |
5–25 |
Strong |
0.6% |
|
F |
6,000–7,500 K |
1.04–1.4 |
1.15–1.4 |
1.5–5 |
Medium |
3% |
|
G |
5,200–6,000 K |
0.8–1.04 |
0.96–1.15 |
0.6–1.5 |
Weak |
7.6% |
|
K |
3,700–5,200 K |
0.45–0.8 |
0.7–0.96 |
0.08–0.6 |
Very weak |
12.1% |
|
M |
2,400–3,700 K |
0.08–0.45 |
≤ 0.7 |
≤ 0.08 |
Very weak |
76.45% |
Giant and Super Giant stars are so large that the densities in their outer regions are low, which subtly changes the appearance of the stellar spectrum. Each spectral class in fact has its own set of criteria. Therefore we can tell if a star is a giant, super giant, or of any other category, only from studying its spectrum.
It is now clear that stars can have different combinations of luminosity and temperature. To identify the combined property (essentially the location in the HR-diagram) a luminosity classification has been developed, the Morgan-Keenan luminosity classification. Roman numerals are used to indicate their MK class, "I" for super giants, "II" for bright giants, "III" for giants, "IV" for "sub giants" (stars that are developing into giants), and "V" for the main sequence.
Examples:
- Vega is an A0 V star,
- Polaris is F7 I or II, and
- Aldebaran is K5 III
- The Sun is a G2 V star.
White Dwarfs are just called White Dwarfs, or class D. All these classes show up distinctly on the HR-Diagram. This diagram has now been improved very much with the help of additional information about typical stars, such as distance and luminosity, and this diagram now plays a very important role in modern astronomy.
It has become an essential tool for interpreting the "finger prints" of stars and reveals a great deal of information, no matter how far away the sources are.
Using the HR-diagram
We discussed that we can find out various properties of stars simply from their position in the HR-diagram. Let us now discuss a few typical examples related to the size of stars.
If we compare all stars that fall in the spectral M-class (figure left) we see at the bottom Proxima Centauri and Barnard's Star. These are red dwarfs that fuse hydrogen into helium but do so very slowly because they are relatively small.
Going up in the diagram we find a star called Mira which is much more luminous than the red dwarfs. Further up we find the very luminous Antares and Betelgeuse.
Luminosity here depends very much on the size of the star, hence its surface area. Low temperature stars that are much larger than any other, will emit more energy per unit of time, i.e. have higher luminosity.
If we now do the same for the hotter spectral A-class, we see at the bottom the White Dwarf Procyon B which is very hot but also very small, i.e. has relatively small surface area and thus low luminosity.
The main sequence star Sirius has a similar temperature but is much bigger, and is actually a main sequence star.
At the top we see the extremely luminous stars Rigel and Deneb, which shows that these stars are much bigger than Sirius and classified as Super Giants.
Finally we compare stars with similar low luminosity. The White Dwarf Procyon B has the same low luminosity as e.g. Barnard’s Star, but is much hotter. This means that while they have similar luminosity, Procyon B must be much smaller (has less surface area) than Barnard’s star.
Finding Mass
As we saw above, for main sequence stars we can infer the mass of the star from where it is in the main sequence, or from its spectral class. Because all these stars fuse hydrogen into helium, the main process during a stars’ lifetime, the difference in spectral class or temperature must be directly related to the mass of the star. Stars with more mass have a higher core temperature and thus fuse hydrogen faster than low mass stars. They also have a shorter lifetime. We discuss this further in our EBook ‘Stellar Evolution’.
Line spectra outside the visible spectrum
Uv- and X-ray emission lines
Absorption and emission spectra also occur in the UV and X-ray parts of the spectrum. That was only discovered since astronomers can launch telescopes into space to observe these parts of the spectrum. These absorption spectra come from the outer atmosphere of the star, the corona and the region between photosphere and corona, where the temperatures are much higher than at the photosphere of the star.
Above image shows X-ray emission lines of the massive star Zeta Orionis (or Alnitak, the faintest of Orion’s Belt) observed by the spectrograph of the Chandra satellite. It is the brightest O-type star with an estimated 31,000 degree K surface temperature. Credit: chandra.harvard.edu.
Thus electrons can be exited much more and consequently their energy jumps can be larger, emitting or absorbing photons with much higher energy (shorter wavelengths). As an example the typical photosphere temperature of our Sun is 6000 K while the corona temperature can go up to several million degrees K. So with observations from space we now can also study the stellar corona from line spectra.
Radio emission lines
Since the early 1950’s radio astronomers observing the Radio part of the spectrum with their radio telescopes, predicted spectral lines in the radio spectrum. Because the wavelengths here are much longer, the energy that causes such spectral absorption or emission must be much smaller than that occurring with electrons that change their energy state.
Visible (inverted) light (left) and 21 cm Radio (right) image of the M81 galaxy system observed by the VLA radio telescopes. The radio image clearly shows the interaction between the various components of this galaxy system. Image courtesy of NRAO/AUI at http://images.nrao.edu
But another process inside the Hydrogen atom, the so-called spin-flip transition releases a much smaller amount of energy and causes a line at the 21 cm wavelength. This hydrogen line was predicted in the 1940’s by the Dutch astronomers Oort and Van de Hulst and has indeed been found with radio telescopes. These observations have been extremely helpful for the study of the Milky Way and other galaxies.
The advantage of radio waves is that they penetrate dust lanes that obstruct us from seeing large parts of our galaxy in the visible spectrum. Therefore this radio technique based on the 21 cm line has been vital for our understanding of the structure and behaviour of our Milky Way and other galaxies.
Molecular Spectroscopy
This can be used to detect molecules that consist of several atoms. Now many compounds have been found, especially in inter stellar space inside galaxies, and we now know that such matter is very diverse and common in the Universe. The combination of spectral research in different wavelength regions from infrared to microwave continues to be of great importance for understanding the structure of galaxies and the makeup of inter stellar space.
Conclusion
In examining spectra, astronomers have developed powerful techniques to determine a lot of detail about stars, galaxies and composition of inter stellar space, in spite of their great distances.
We can:
- determine chemical composition and temperature,
- distinguish stellar classes from sub-dwarfs all the way to super-giants,
- and thus estimate mass, age and luminosity.
In other EBooks we can learn more about how these spectra help us to determine distance and radial velocity. With that information we can study multiple star systems, find planets around other stars and even observe the expansion of the Universe.
Credit: Menlo Systems
The great advantage of spectral astronomy is that it does not matter how far away the celestial object is. Astronomers even study the spectra of objects in the outer reaches of the known Universe at distances up to 13 billion light years.
There is continued R&D in the technique of spectroscopy, which paves the way for ever more exciting discoveries. See e.g. here
Further reading
- In the EBook “Stellar Distance” we further discuss the link between luminosity and distance and we also explain how the line spectrum is used to very accurately measure radial velocity in space;
- In the EBook “Stellar Evolution” we discuss the life cycle of stars and will come back to the HR-diagram.