Chapter 7: Projects and Measurements

 

In this chapter, I present several projects that I have done with the SGS. These are merely sample projects that could be developed into full projects. I used them as a means of learning how to perform spectrometry in general.

 

7.1 Spectral Classes

In the early 1900’s, astronomers and physicists began to understand how the temperature and density of a star would affect the features in its spectrum. In the 1920’s a classification system was designed that classified stars by their temperature as determined from their spectral profiles (Kaufman 468). There are seven major groups in the core classification system, designated by the letters O B A F G K M. The hottest stars are the O class and have temperatures over 25,000 K. At the other end of the series are cool M stars, roughly 3,000 K. These spectral classes are broken down further into spectral types. The spectral type is defined by a number after the class letter. For example, the star Rigel is a B8 star. The spectral types were added to the classification for finer detail. Not all stellar classes have an equal division into spectral types. The B class for example has more subdivisions than the G class. Table 7.1 has the full range of spectral classes in use today.

 

Table 7.1Kitchin 190

O4

B0

A0

F0

G0

K0

M0

O5

B0.5

A2

F2

G2

K2

M1

O6

B1

A3

F3

G5

K3

M2

O7

B2

A5

F5

G8

K4

M3

O8

B3

A7

F7

 

K5

M4

O9

B5

 

F8

 

 

M7

O9.5

B7

 

F9

 

 

M8

 

B8

 

 

 

 

M9

 

B9.5

 

 

 

 

 

Stellar spectra are marked by absorption lines caused by elements in the stars' atmosphere. The intensity of the absorption lines depends on the balance between ionization of these elements and the excitation of their electrons to certain levels. For example, hydrogen Balmer lines are most prominent in the A0 to A5 spectral types that are around 10,000 to 7,500 K. If the star is hotter than this, it emits radiation at levels high enough to begin ionizing the hydrogen. The hotter the star becomes, the more hydrogen is ionized and the fewer the number of hydrogen atoms available to be excited to the Balmer energy levels. Fewer atoms being excited produces a weaker absorption line. For stars below the 7,500 K, not enough of the emitted radiation is strong enough to excite the atom to these energy states and the intensity of the absorption line decreases (Kaufman 469). Thermodynamics allows astronomers to accurately determine the temperature of the star based on the intensities of the absorption lines.

The spectral type is determined by the relative intensities of two carefully chosen absorption lines. If one absorption line increases with temperature while another decreases, the ratio of the two can be used to accurately determine the temperature. These variations can be distinguished by the naked eye; however, Vspec allows us to measure intensities quantitatively, making this process easier. The line pairs are different for different classes, and unfortunately, they are usually not the dominant lines in the profile. The more sensitive lines generally fall in the blue-green region, 3500 to 5000 angstroms, and tend to be weaker lines. For example, the temperature of A stars is determined by the ratio of Mg II (4481Å) and Fe I (4385Å). These are not as prominent as the hydrogen Balmer lines that dominate the profile (Kitchen 193 195).

Figure 7.1 shows the profile of several bright stars ranging through the entire spectral series. These spectra were taken with the narrow slit and low-resolution grating. The exposures were 1 to 2 minutes and unguided. Their profiles were processed against a standard spectrum as described in chapter 6, and I have labeled some of the prominent lines. Notice how these lines change with temperature, especially the hydrogen lines. Also, notice that the continuum moves from a peak in the blue to a peak in the red as the temperature decreases, accurately following the black body radiation curve.

The stars used to create this list were, in order from hottest to coolest:

 

Mintaka

Spica

Rigel

Vega

Seginus

Wasat

Procyon

Sun

Capella

Asellus Austrailis

Arcturus

Aldeberan

Betelguese

 

There is an artifact in these spectra left over from the continuum modeling performed in Vspec. In the red end of the spectrum there is an absorption line produced by O2 in Earths atmosphere. The altitude of the star when its spectrum was taken determines the intensity of this line. At low altitudes, the light from the star passes through more of Earth's atmosphere and therefore more light is absorbed. However, in the case of these spectra, the noticeable trend is greater intensity for cooler stars. This is not because I took the spectra of the cooler stars while they were low on the horizon. Rather, the continuum correction in Vspec weakens the red end of the spectrum when compensating for blue loss in the spectra of hot stars. Thus, the strength of this atmospheric line is severely weakened for hot stars in comparison to cool stars.

Spectral classes have been determined for most of the stars that are attainable with our scopes, so this is not ground breaking work. Still, it can be a fun and is good exercise to start with.

Figure 7.1

Format adapted from Kaufman 47.

 

 

7.2 Intensity Classes

Stars are also classified by their luminosity. Stars of the same temperature and the same spectral class can be different sizes. The size of a star affects its luminosity; larger stars are more luminous than smaller stars. The luminosity does not affect the overall profile of the star; however, the size of the absorption lines changes slightly. The atmospheres of smaller stars are denser and atoms in the atmosphere collide more frequently. When atoms collide they distribute some of their energy to each other. This causes the atoms to absorb a wider range of light frequencies, causing the absorption lines to be broader. This effect is known as collision broadening. Astronomers can determine the luminosity of a star by the size the absorption lines. Luminosity class is designated by a roman numeral after the spectral classification.

Astronomers measure the amount of energy removed or absorbed from the continuum by an absorption line’s equivalent width. The equivalent width is the width of a rectangular line that has the same area as the spectral line. This allows astronomers to easily compare the intensity of two lines with different shapes. In Figure 7.2, the rectangular line superimposed on an absorption line.

 

Figure 7.2

Adapted from Kaufmann 472

 

 

Figure 7.3 shows profiles of Rigel and Algol, both are B8 stars, but Rigel is a blue super giant and Algol is a smaller main sequence star. Identified are two hydrogen Balmer lines. I used Vspec to measure the equivalent width of each line. The widths are given beside each line; the negative values refer to the absorption. Notice that as predicted the smaller star has wider absorption lines.

 

Figure 7.3

 

7.3 Atmospheric Lines

Our own atmosphere has molecules that absorb light. There are three dominant absorption lines in our atmosphere. Two are from the oxygen molecule, O2, and are located at 6875 and 7590 angstroms. Water in the atmosphere absorbs light between 6950 and 7250 angstroms. These lines can be seen against the continuum of a star. Their strength depends on the altitude in the sky at which the spectra was taken. The closer to the horizon, the more atmosphere light has to pass through and therefore there is more atmospheric absorption.

Early on in my research, the spectrometer was calibrated incorrectly and I was examining the near infrared rather than the visible spectrum. I mistook the O2 line at 7590 for the Hydrogen alpha line. I was perplexed as to why the intensity of this line did not change between spectral classes. It is this type of set back that I hope my thesis will prevent other students from running into.

Figure 7.4 shows the spectra that I confused. The 7590 O2 line is the dark band on the far right of the spectra below. The contrast of the spectra has been adjusted to make it easier to see this line; as a result, the rest of the spectra appear saturated.

 

Figure 7.4

 

Figure 7.5 shows the profiles of three different stars. The profile is much more helpful for identifying lines. The double hump of the O2 lines immediately identifies them; this shape cannot be easily seen in the image itself.

 

Figure 7.6

 

7. 4 Methane Absorption

The atmospheres of planets absorb light just like the atmosphere of stars. However, because the atmospheres of planets are much cooler, larger molecules, rather than individual atoms, are more likely to absorb light. These absorptions generally occur in the infrared spectrum. Fortunately, our camera is sensitive to near infrared light. I was able to observe methane absorption bands in Jupiter and Saturn, see Figure 7.7 and Figure 7.8. Notice that methane is not present in Saturn’s rings. Light at those wavelengths is absorbed by the methane on Saturn is reflected by its rings.

 

Figure 7.7: Jupiter

Figure 7.8: Saturn

 

 

 

7.5 The Orion Nebula

The spectrum of the Orion Nebula in Figure 7.9 is an example of a two-dimensional spectrum. The density and temperature of the nebula are not consistent, and these two properties greatly influence the strength of the nebula’s emission lines. In this high-resolution spectrum, there are three prominent lines. From right to left they are 6548.09 from ionized nitrogen, NII, 6562.852 which is the H alpha line, and 6583.36, also from NII. Notice that they change intensity at different vertical positions, indicating a change in temperature or pressure in the nebula. The two NII lines are useful for determining the temperature of the nebula, the ratio of the sum of their intensities to the intensity of the 5755 NII line can be used to find the temperature (Osterbrock 102). Astronomers in the past have sometimes had difficulty trying to resolve these lines. Our high-resolution grating does it easily. The only drawback is that our spectral range is too narrow to get the 5755 lines in the same image.

 

Figure 7.9

 

Although we cannot determine the temperature of the Orion Nebula from this image, we can use the image to prove some quantum mechanical theory. The ratio of the transition probability for the two transitions producing the NII lines is exactly 3:1 (Osterbrock 50). In order to analyze the spectrum I had to isolate certain sections. In Figure 7.10, I divided the region into three sections, A, B, and C. I analyzed their spectra with Vspec and measured the equivalent width of each line. I found the results astonishingly precise; they are shown in Chart 7.1.

 

Figure 7.10

 

 

Chart 7.1


 


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