Universe's
finger prints – Spectroscopy
Remote
telescope spectroscope monitoring!
We know what the stars are made of, know of their structures and
their lives, only because we are able to observe and analyze their spectra.
Unbroken starlight allows us to admire a star's external characteristics; its
spectrum allows us to look into its very soul.
More than 300 years ago Sir Isaac Newton (1642-1727) showed
that sunlight can be splitted into different colors
using a simple prism.
He found that the
shorter the wavelength the greater the angle of refraction so that a spectrum
of light is produced from red through to violet.
Stellar Spectroscopy is the study of the spectra of starlight.
It is a very powerful tool that enables astrophysicists to
infer many physical and chemical properties of stars and classify them into a
logical sequence.
In order to understand how spectroscopy can be a very useful and meaningful tool
to astrophysicist one need to understand the different kinds of spectra that
are observed.
Pic: The electromagnetic spectrum. From Gamma to Radio.
Spectroscopy is the study of the characteristic
wavelengths or colors. Optical emission spectroscopy comprises several
techniques that form the most important means we have for chemical analysis.
In spectroscopy we measure spectra emitted by
atoms and ions with optical transitions in the wavelength range from about 400 nm
to 850 nm. This range includes the ultraviolet, and visible light (from
violet at 380 nm to red at 760 nm), and the near infra-red.
·
Spectroscopy rule of thumbs (Kirchoff's):
Rule 1
A hot opaque solid, liquid or gas which is under high pressure will emit
a continuous spectrum.
Rule 2
A hot gas under low pressure (much less than
atmospheric) will emit a series of bright lines on a dark background. Such a
spectrum is called a bright line or emission spectrum.
Rule 3
When light from a source that has a continuous
spectrum is shone through a gas at a lower temperature and pressure, the
continuous spectrum will be observed to have a series of dark lines
superimposed on it. This kind of spectrum is known as a dark line or absorption
spectrum.
Pic: a stars spectrum's representing.
Continuous
Spectra
In a very hot gas, the atoms have high kinetic
energies and collisions between them are very frequent. Their electrons are
raised to a very excited states and then drop down producing emission lines.
However, if the gas is at very high pressure and density, then an electron in
its excited state may not have enough time to drop down to its ground state before
it undergoes another collision from a neighboring atom. This has the effect of
blurring the sharpness of each emission line into a broad band of wavelengths.
The same thing happens to neighboring lines so that by the time the light
emerges from the gas it has 'smeared out' into a continuous spectrum at all
wavelengths.
Emission
Spectra
In a gas containing only atoms of one kind, the
electrons will all be in their ground state if the temperature is low. As the
gas is heated, its atoms gain kinetic energy and collide with their neighbors
causing their electrons to be raised to excited states. As the electrons drop
down, photons will be emitted with many different energies and wavelengths
corresponding to the particular electron energy level scheme for the gas. The
emission of these lines will cause the gas to glow with a light composed of
wavelengths that correspond to the electron energy transitions. For moderate
temperatures we might find that only the first excited state of the atom is
attained and so the emission light will consist of a single bright emission
line corresponding to the difference in energies between the first excited and
ground states. As the temperature is increased, more emission lines will start
to appear - until
at higher temperatures many lines will be visible corresponding to all the
allowed energy transitions of electrons in the gas. In this way an emission
line spectrum is formed that is related to the elements composition of the gas.
Absorption
Spectra
Light from the continuous
source contains photons of all energies and wavelengths. If it is the case that
the energy of some of these photons is exactly equal to the difference between
the ground state and an excited state of an atom in the unknown gas, then that
photon will be removed from the incident light. The excited electron will quickly
return to the ground state emitting a photon however, the emitted photon need
not be emitting along the same direction as the absorbed photon but is usually
emitted in a different direction. The re-emitted photons are not
therefore, generally observed through a spectroscope at the source, and the
continuous spectrum is observed when looking to have dark lines at the
wavelengths corresponding to excited states of the atoms in the unknown gas.
It follows that it is precisely
these wavelengths at which light would be emitted in an emission
spectrum if the unknown gas was heated to a high temperature.
Both the dark lines superimposed on the continuous spectrum and the bright
lines in the emission spectrum provide a 'spectral fingerprint' that identifies
the elements present in a hot gas.
Spectral Types
The spectral type of a star is designated by one of seven
letters O, B, A, F, G, K, M, starting with the hottest
type (O type) to the coolest type (M-type). The table below shows the
temperatures and characteristic features in the star's spectrum
that distinguish spectral types.
A nice method of remembering the spectral order is to put
it in a nice verbal structure : "Oh, Be A Fine,
Girl –
|
Star
Type |
Surface
Temperature / K |
Spectral
Type |
|
O |
>20 000 |
ionized
helium (He II) . |
|
B |
20,000 - 10 000 |
neutral
helium, hydrogen lines start to appear. |
|
A |
10,000 - 7000 |
strong
neutral hydrogen (Ballmer lines) visible . |
|
F |
7,000 - 6000 |
ionized
calcium (Ca II) visible, hydrogen lines weaker . |
|
G |
6,000 - 5000 |
ionized Ca II very prominent,
much weaker neutral H lines, also other metallic lines such as Iron
(the sun is a G-type star) . |
|
K |
5000 - 3500 |
neutral
metals such as Ca and Fe prominent, molecular bands visible. |
|
M |
3500 - 2000 |
molecular
bands very visible, particularly those of Titanium Oxide (TiO)
. |
The wavelength of a spectral
line is affected by the relative motion of the star and the observer.
Due to the Doppler effect - light from a star
will be shifted to the blue end of the visible spectrum if it is approaching
the observer and shifted to the red end if it is receding.
The Spectrograph
Astronomers produce spectra by
means of a "spectrograph" affixed to the telescope. The oldest form
of the device was visual (a spectroscope), and consisted of little more than a
prism in a tube fixed to the end of the telescope, the refracted light focused
by an observer's eyepiece. By the turn of the 20th century, spectra
was recorded with films.
In the middle of the century,
prisms were replaced by diffraction gratings – which are finely ruled surfaces
that produce spectra by the interference of light waves.
In the modern spectrograph,
light is sent from the telescope onto a "collimator" or a "Grating".
The grating makes a spectrum, whereas the colored
light then focused by a camera onto the CCD camera that records the spectra
digitally into the PC. Spectra are commonly seen reproduced either
photographically or graphically (see diagram bellow).
The spectrograph, fitted to the
base of the telescope, breaks the light into its component colors and records
the spectrum.
Bareket observatory astrograph
has about 25A per pixel resolution. Although this is consider
to be a pretty low spectra resolution, it is also very bright;
enabling to
produce a very positive and permanent results from a very wide range of different subjects –
·
As super novas.
·
Quasars.
·
Asteroids.
·
Comets.
·
Stars.
·
Nebulae.
·
And even the planets atmospheres
and moons!
* There is a possibility of
making higher resolution spectra on a special events and / applications upon
request.
