Results and Discussion

The low-resolution mode is useful for stellar classification and obtaining spectra of planetary nebula. In the high-resolution mode, many absorption lines are visible of atoms, ions and simple molecules. Figure 11.3 shows the higher-resolu-

Figure 11.3. Classification of stars based on their spectra. Spectra were obtained in the high-resolution mode from Hp to HS for stars from early B to M class. Note the increase followed by decrease in the hydrogen absorption lines as one proceeds from B- to M-type stars along with the general increase in the number of metal lines as the temperature decreases toward M-class stars. Several different metal lines are identified, along with the G band representing the diatomic molecule CH.

Figure 11.3. Classification of stars based on their spectra. Spectra were obtained in the high-resolution mode from Hp to HS for stars from early B to M class. Note the increase followed by decrease in the hydrogen absorption lines as one proceeds from B- to M-type stars along with the general increase in the number of metal lines as the temperature decreases toward M-class stars. Several different metal lines are identified, along with the G band representing the diatomic molecule CH.

tion spectra of stars from class B to M and luminosity class III and spans the region from Hp to about 4100 A. Several of the more prominent lines are labeled. As one can see, many lines are present, especially as one proceeds to cooler stars. Graphic display of these spectra, for example, using Excel, results in line profiles, especially for the more intense lines such as those for hydrogen. These profiles contain information regarding the physics of the stellar atmosphere. For example, buried in the line profile and intensity are such information as pressure, density, abundances and temperature, some of which, with the appropriate mathematics or software tools, can be extracted. In addition, rotation of the star is also contained within the profile, which also can be extracted [4].

Figure 11.4 shows the emission line spectrum of NGC7009, the Saturn Nebula, in the lower-resolution mode. The upper-left corner shows an image of the nebula with the slit superimposed over the nebula. The main body of the line spectrum was obtained in the low-resolution mode so that a majority of the emission features could be observed in a single spectrum. A higher-resolution spectrum is shown in the upper right, centered at the Ha line. In high-resolution mode, three components are seen to make up this strong emission, Ha and two lines of N+, which flank Ha. A graphical presentation of the line spectra is also shown. The usual types of species are apparent in the spectrum of a planetary nebula. The hydrogen and helium lines, along with the prominent forbidden O+2 lines, Ar+2, Ar+3, N+2 and S+, round out the variety of ionic species present. For extended objects, such as nebulae, one obtains a spectrum across the entire length of the slit. During the analysis of the results, the software allows you to select small subregions of the spectrum for analysis. As a result, one can profile the composition across the entire slit, effectively obtaining many spectra of the object at a single time, which is analyzed within the software. As discussed later, this would also allow for profiling temperature and density parameters of the nebula across the slit.

In addition to identifying emission features, other physical features of a nebula can be determined, such as the temperature (T) and electron density (Ne). Theoretically derived equations [5] have been obtained that relate electron density and temperature to line ratio intensities, as seen in Equations (1) and (2).

Equation (1) relates the line-intensity ratios for transition occurring at 4959, 5007 and 4363 Angstroms for O+2, while Equation (2) relates the lines at 6548, 6584 and 5755 Angstroms for N+. The bottom panel of Figure 11.4 shows the graphic solution of these equations, yielding a temperature and electron density. This example portrays the gold mine of information contained within a spectrum. As was mentioned earlier, for an extended object the spectrum across the entire slit is obtained and line ratios as a function of distance from the central star, for example, can be obtained in a single spectrum. The potential is present to create two-dimensional profiles of temperature, density and composition for nebular-type objects.

When the spectrometer is used in high-resolution mode, many absorption features can be observed in the spectra, particularly in cooler stars. Simple image-processing techniques enhance these features making their identification much easier. Figure 11.5 shows the high-resolution spectrum of 13-theta Cygnus, an

(¡4959 + I5007)/I4363 = [7.15/(1 + 0.00028Ne/T1/2)]10 (I6548 + ¡6584)/¡5755 = [8.50/(1 + 0.00290Ne/T1/2)]10

,14300/T

,10800/T

Figure 11.4. Spectrum of planetary nebula NGC 7009 (Saturn Nebula). In the upper-left corner of the top panel, the positioning of the slit is indicated. The low-resolution spectrum is shown as both a graph and an emission line profile. The high-resolution spectrum is shown in the upper right, centered at the Ha line showing the presence of N+ lines flanking the Ha line. Various other ionic and atomic species are identified. Many of these lines are forbidden, such as the intense 5007 Angstrom line of O+2. The bottom panel shows the kind of physics that can be derived from such a spectrum. Theoretical equations describe the relationships between electron density and temperature to line intensity ratios. Two such equations are shown (there are many). The ratios are determined from the calibrated spectra and the two unknowns solved to yield electron density and temperature.

Figure 11.4. Spectrum of planetary nebula NGC 7009 (Saturn Nebula). In the upper-left corner of the top panel, the positioning of the slit is indicated. The low-resolution spectrum is shown as both a graph and an emission line profile. The high-resolution spectrum is shown in the upper right, centered at the Ha line showing the presence of N+ lines flanking the Ha line. Various other ionic and atomic species are identified. Many of these lines are forbidden, such as the intense 5007 Angstrom line of O+2. The bottom panel shows the kind of physics that can be derived from such a spectrum. Theoretical equations describe the relationships between electron density and temperature to line intensity ratios. Two such equations are shown (there are many). The ratios are determined from the calibrated spectra and the two unknowns solved to yield electron density and temperature.

F4-type star, demonstrating the variety of elements, which can easily be identified. A few of the many lines have been identified and are labeled in the spectrum. The upper spectrum of the top panel identifies unionized elements in the spectrum, which spans the blue part of the visible spectrum from Hp to HS, The lower spectrum of the top panel identifies ionized elements and covers the

Figure 11.5. The upper panel shows the spectrum of 13-theta Cygnus and an F-type star in the blue region. The upper spectrum of that panel identifies a variety of atomic species, while the lower panel shows the same spectrum but identifies ionic species. The presence of absorption lines, mostly of metals, is even more pronounced in cooler M-type stars as shown in the lower panel. This is a spectrum of omicron Ceti (Mira) near maximum light and shows the great number of lines, many of which are blends of several different lines. This spectrum spans only 160 Angstroms and is the result of using a higher-resolution grating than the one used in the upper panel.

Figure 11.5. The upper panel shows the spectrum of 13-theta Cygnus and an F-type star in the blue region. The upper spectrum of that panel identifies a variety of atomic species, while the lower panel shows the same spectrum but identifies ionic species. The presence of absorption lines, mostly of metals, is even more pronounced in cooler M-type stars as shown in the lower panel. This is a spectrum of omicron Ceti (Mira) near maximum light and shows the great number of lines, many of which are blends of several different lines. This spectrum spans only 160 Angstroms and is the result of using a higher-resolution grating than the one used in the upper panel.

same wavelength region. Even more dramatic are the lines present in cooler stars such as are seen in the lower panel of Figure 11.5, which is a spectrum in higher resolution (0.5 Angstrom/pixel) of omicron Ceti (Mira), an M-type star. At this point let us diverge a bit and talk about lines.

The identification of species responsible for the observed absorption lines remains somewhat of an art, since no software packages are available that will do this. The Chemistry and Physics Handbook lists more than 21,000 lines between wavelengths of 3600 and 10,000 Angstroms for all the elements. This represents, on average, about three lines per Angstrom wavelength interval. Since the instrument is only at best able to resolve 0.3-0.5 Angstrom in the highest resolution mode, some criteria must be used to eliminate many of the potential lines observed or, as is often the case, a line may represent a blend of two or more lines. The presence and intensity of a feature, due to an atom or ionic species, is the result of many parameters such as natural abundance, probability of the absorption, temperature, density and pressure. For our purposes, the first three are the most important. Many potential features can be eliminated at the start simply because the natural abundance of an element is so low. For example, at a given temperature the Hp absorption line at 4861.33 Angstroms has a relative intensity of 80 compared to 3000 for protactinium at 4861.49 Angstroms. Yet one would not typically assign this absorption line in a stellar spectrum to protactinium, simply because the natural abundance of this element is 9 orders of magnitude lower than that of hydrogen. In addition to this abundance factor, other absorption lines for hydrogen are present where they should be, whereas protactinium lines are not. To be certain of assignment of a line, it is important to find several lines for a particular species. In addition, for most of the spectra shown in this chapter, the spectra obtained were compared directly with those obtained in the astronomical literature to be certain of the assignments. My initial goals in using this instrument were to establish how far the instrument could be pushed in giving spectra and in the resulting identification of features.

Figure 11.6 illustrates just how far this instrument is capable of being pushed to give useful data. 19 Piscium is a carbon-type Mira variable star. These type of stars are my particular interest and represent old, evolving stars nearing their deaths. Stars of this spectral class often exhibit absorption lines owing to the unstable element technetium. This element is not found naturally in the solar system because its longest lived isotope 97Te, has a half-life of only 2.6 x 106 years and, as a result, all technetium endogenous to the formation of the solar system has long since decayed. Yet multiple lines of technetium can be detected in the atmosphere of this star and others of the S and C types [6]. It is known that at certain stages in the evolution of these stars neutron capture reactions, on seed iron-nickel atoms, are proceeding deep within stars of this type. At other stages in the stars' evolution, dredging mechanisms bring to the surface these nuclear-processed materials, now well laced with heavy elements including technetium. These heavy elements such as zirconium and zirconium oxide molecules, now often appear in the stars' atmospheres, which give rise to extensive banded structure in the spectrum in the red region.

The other very interesting aspect of this type of star is the fact that they often contain abnormal amounts of 13C compared to 12C. The normal solar system ratio of 12C to 13C is 80, but in many of these type of stars this ratio can approach 4. This follows from theoretical work, which suggests that the neutron source in these stars, which gives rise to neutron capture heavy elements, is due to the burning of pockets of 13C deep down in the star. The dredge-up results in the presence of these heavy elements in the stellar atmosphere and can result in

Figure 11.6. The presence of the unstable element technetium is shown in the carbon star 19 Piscium. Many stars of the S and C types exhibit these lines and indicate that heavy element synthesis, via slow neutron capture processes, are occurring in these evolved stars and being dredged up to appear in the surface layers. Since technetium is an unstable element with a half-life of only a few million years, this indicates that heavy element synthesis is an ongoing process, not limited to supernova explosions. Many other heavy elements are often also visible such as strontium and zirconium.

Figure 11.6. The presence of the unstable element technetium is shown in the carbon star 19 Piscium. Many stars of the S and C types exhibit these lines and indicate that heavy element synthesis, via slow neutron capture processes, are occurring in these evolved stars and being dredged up to appear in the surface layers. Since technetium is an unstable element with a half-life of only a few million years, this indicates that heavy element synthesis is an ongoing process, not limited to supernova explosions. Many other heavy elements are often also visible such as strontium and zirconium.

enhanced quantities of 13C. Normally, the detection of isotopes of atoms or ions cannot be done easily because the lines are extremely close together and normal line-broadening effects cause them to overlap. However, this is not the case with molecules where rotation and vibration transitions are much more sensitive to the isotope composition. These types of transitions are normally outside the optical range, but electronic transitions can couple with vibration transitions giving rise to a blanketing effect of the multitude of lines that result, and these lines are often observable in the optical region. In cooler S-and C-type stars, diatomic carbon forms and a clear separation of absorption lines occurs between diatomic 12C-12C, 12C-13C and 13C-13C as indicated on the spectrum in Figure 11.7.

Figure 11.7 shows the graphical spectra of three different stars containing varying ratios of 12C to 13C. Some of these stars are noted for their large quantities of carbon as observed with diatomic carbon. All three possible combinations of diatomic C-C (soot) are observed, 12C-12C, 12C-13C and 13C-13C. Spectrum A represents a carbon star with a solar system-like ratio of the two isotopes (12C/13C of around 80). In spectrum B this ratio has decreased to 20 and results in the appearance of 12C-13C in the spectrum. Finally, in spectrum C the ratio is around 4, and 13C-13C makes its presence known in the spectrum. The blanketing effect of diatomic carbon is shown along with the identification of several metal lines.

In addition, with careful wavelength calibration, one can measure the Doppler shift of absorption and emission lines to determine velocities of approach or recession of objects along with rotation velocities of stars and planets. I have been w C

»C

t

, A

_ /V A f^J ^

MiJti'idm t".

- A

u y r tan J

ü

Ii Vt

' A/

c

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Wavelength (A)

4780

4720 4740 4760

Wavelength (A)

4780

4800

Figure 11.7. The spectra of three different carbon stars (C stars) are shown in this graph in the 4700-4800 Angstrom region of the spectrum. Proceeding from A to B to C, the 12C/13C ratio decreases from 80 to 20 to 4. This results in the gradual appearance of diatomic carbon (C2) containing various combinations of carbon isotopes. Note the large blanketing effect C2 has on the spectrum blueward of 4740 Angstroms. This can be so dramatic in these type of stars that the blue part of the spectrum can essentially disappear.

able to successfully observe and measure the rotation of Saturn by aligning the slit along the rings of the planet. In only a few-second-long exposure one is able to obtain a spectrum of the entire disk and ring system of Saturn. The eastern and western region of the image can be isolated using SPECTRA software, and each of the spectra calibrated against the Ha line. When the final two images are aligned, they show the clear shift in the lines, which occurs due to Saturn's rotation. From this the rotation velocity can be calculated. In a similar manner, the velocity of approach of M31 has been determined by acquiring the spectrum of the nucleus of the galaxy. In this case, only the narrow 18-|im slit was used, and as a result a relatively long 60-minute exposure was necessary. I have not tried using the wider slit, but according to the instrument specifications, a gain of 2 magnitudes is possible, with a loss of resolution, which would occur with the wider slit, but one should be able to determine the center of lines with the software provided.

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