Discussion
Assignment 5: Principles of Spectroscopy
Due:
Beginning of discussion on Thursday, October 20
Name:
Section (Circle one): 004 005
Introduction
At the center of the sun, hydrogen nuclei collide and combine to create helium. In doing so, energy is released -- energy that eventually comes to us in the form of light. This is about the journey of each individual 'piece' of light, or photon, from the center of the sun to us, and the information that we gain as a result.
A photon is a small bit of electromagnetic energy sent across space. Photons can be emitted or absorbed by electric charges -- usually an electron.
A hot, dense object contains many "loose" electrons which can emit photons of any energy. However an electron cannot emit a photon with more energy than the electron started with. The light produced by a hot, dense object is called continuous emission because it contains photons of all energies, i.e. light of all colors, or wavelengths. The resulting "rainbow" is called a continuous spectrum. As we heat up an object, we are giving the electrons more kinetic energy, so they become able to emit more energy. The hotter the object becomes,the brighter the continuous emission becomes. This is described by the Stephan-Boltzmann Law:
f = σT4
As the emitting object is heated, the flux, f, of light energy emitted per unit area (the brightness) increases as the temperature, T (measured in Kelvin, K), to the fourth power; σ is called the Stefan-Boltzmann constant, and has the value 5.67x10-8J m-2 K-4. If two hot pokers are the same size, but one is twice as hot as the other, the hotter one will be sixteen times brighter. The same is true of two stars.
As the object heats up and the electrons get more energy, the energy of the typical photon emitted also increases. This means that the continuous emission gradually shifts toward shorter wavelengths (higher energies) and therefore looks bluer. This is described by Wien's Law:
λpeak * T = 0.29 [cm K]
which means that as the temperature, T, of the emitting object increases, the wavelength λpeak where the intensity of the light is the greatest must decrease. A very hot poker will glow with a bluer (shorter wavelength) light while a cooler poker will glow with a redder light.
Any hot, dense, opaque object can and must produce continuous emission across all wavelengths, with the total energy and dominant color described by these two laws. This is sometimes called blackbody radiation or thermal radiation. The object has no choice -- if it's hot, the electrons have energy, so they must emit light. Remember, Wien's law and the Stefan-Boltzmann Law apply only to continuous thermal emission.
So far we've talked about processes involving "loose" electrons that lead to thermal radiation. What about electrons that are part of an atom? In the Bohr model of the atom, electrons orbit a nucleus of protons and neutrons. Each orbit has a different potential energy, just like planetary orbits correspond to particular gravitational potential energies. But according to quantum mechanics, the electrons can only orbit in certain places, which means the electrons can only have certain orbital energies -- these allowed energies are called energy levels.
Electrons usually stay in low energy levels, but they can "jump up" to higher energy levels by absorbing a photon or by gaining energy in other ways. If it gains energy by absorbing a photon, it has to have exactly the correct amount of energy -- it has to match the energy difference between the energy levels. Therefore, the atom can only absorb light at a few specific energies, or colors. This is called line absorption. Line absorption occurs when a low-density gas is in front of a hotter, continuous emission source. The cooler, low-density gas acts to block the photons which have the right wavelengths, while the other photons travel through the gas unperturbed. This leads to a generally bright spectrum, with dark lines at specific wavelengths. The missing colors are called spectral absorption lines and result in an absorption spectrum.
The energy-level jumping can also happen in reverse. The electron can "fall down" from a higher energy level to a lower one, emitting a photon with energy equal to the difference between the levels. This is called line emission, because photons are emitted. The spectrum produced is a set of bright emission lines, so it is called an emission spectrum. This can only occur in a low density gas viewed on its own or in front of a cooler background (if a hot, dense object is in the background, we see line absorption instead of line emission).
Notice that these two processes only involve photons with particular energies that match its energy levels. Since each atom or molecule has a different set of energy levels, each atom or molecule also has a unique set of spectral lines.
Let's summarize what are known as "Kirchoff's Laws." First, a hot, dense gas (or a solid or liquid) has free electrons and will emit a continuous spectrum, with the brightness and typical color described by the Stefan-Boltzmann and Wien Laws. Second, a low-density gas along the line of sight to a hotter continuous radiation source will absorb photons of specific energies, leaving an absorption spectrum. Third, a low-density gas viewed alone or in front of a cool background will produce an emission spectrum.
As photons travel outwards from the center of the sun, where the density and temperature are high enough to allow fusion, they are constantly absorbed and re-emitted by the atoms in the sun. Eventually they get to the outer edge of the sun, called the photosphere, which is where the sun changes from being opaque to being transparent. The photospere, then, is the layer where all the photons we see originate. The transparent region above the photosphere is called the atmosphere of the sun and has two major layers. The cooler thin layer abover the photospher is the chromosphere. Above that is the increadably hot and thin Corona.
One photon by itself can't tell us much about the photosphere or atmosphere, but by looking at all the photons together, astronomers can gain information about the temperature, density, and chemical composition of the sun. This is done by looking at the spectrum of the light -- the number of photons (i.e. the brightness) at each wavelength. Similarly, the characteristics of the spectra we will look at in the lab will tell us information about the sources of light we will use.
In the photosphere (and regions deeper in the sun), the density is so high that the gas is opaque. This area produces light with a continuous spectrum. It is radiating simply because it is hot.
First get a dispersion grating from your instructor. A dispersion grating does the same thing as a prism: it splits up the light into individual wavelengths so that you can see the spectrum.
Your instructor will turn on the power supply and the lamp. This dispersion grating is designed to show you many spectra from the same source. Pick a bright, wide spectrum close to the filament that you can tell the colors apart fairly well. Note there may be smaller spectra from reflections or other light sources. Once you've figured out which spectra you want to observe, answer the following questions.
The temperature of the filament can be changed by changing the voltage going to the light bulb. We will start with the light on the lowest voltage necessary to see the spectrum, then slowly turn it all the way up to 120 V. Watch the spectrum as the voltage is turned up, especially the relative strengths of the colors. The brightest color is the peak color and is generally the color of the filament. HOWEVER, our eyes and brains adjust quickly to the light, dimming the brightest colors and reacting poorly if the light is bright (so the filament will never look blue-green). Because of this, you must take your first impression from the spectrum. The best thing to get the peak color is to close your eyes for a few seconds then take your first impression from the spectrum.
Now we will look at "discharge tubes," which are each filled with a low-density gas made of a single kind of atom. Running an electric current through the discharge tube gives the electrons energy and kicks them up to a high energy level. The electrons quickly fall back to their original energy level, giving off a photon with a wavelength determined by the difference in energy between the levels. Most of these photons leave the gas without interacting with other atoms allowing us to view them. This is similar to the process that occurs in the low-density, incredibly hot outer most regions of stars called the corona and in low-density, gas clouds in space called emission nebulae.
There should be a spectroscope set up to observe the spectrum. A slit is aligned with the light source that allows light to travel down to the diffraction grating at the eyepiece. The spectrum is projected onto a scale to the left of the light source. Observe the spectrum through the spectroscope.
element:
element:
Color and mystery number:
| Star | Temperature | λpeak | Color |
| Sirius | 10,000 K | ||
| Sun | 5,800 K | ||
| Betelgeuse | 3,000 K |