Spectroscopic Atlas for Amateur Astronomers 101
TABLE 62 Chi Cyg
S7/1.5e
TiO 7126 R Cygni
TiO 7088 S4/6e
TiO 7054
Richard Walker 2011/03©
Telluric O2
TiO 6815 ZrO 6988
TiO 6715 ZrO 6933
TiO 6681
TiO 6651 ZrO 6778/87
ZrO 6743
TiO 6569
TiO 6478 ZrO 6475/81/94 Hα 6562.82
ZrO 6132/36
TiO 6358 VO 6532
TiO 6180 ZrO 6378/84
TiO 6159 ZrO 6350
Na I 5890/96 ZrO 5839/49
TiO 5847
TiO 5814
TiO 5668 4 Ori ZrO 5718/24
TiO 5597/03 M3.2S ZrO 5629
TiO 5448/51 ZrO 5545/51
ZrO 5404/07
ZrO 5375/79
ZrO 5298/5305 ZrO 5246
TiO 5167
TiO 5003/20 Hβ 4861.33
TiO 4965
ZrO 4792
TiO 4761-63 ZrO 4737
TiO 4626
TiO 4584 ZrO 4641
Sr I 4607
Ca I 4226.73 Ba II4554
ZrO 4471
Hγ 4340.47
Tc I ?
Tc I ?
Hδ 4101.74
Spectroscopic Atlas for Amateur Astronomers 102
HR PegasiTABLE 63
S4/1+
TiO 7126
BD CamelopardalisTiO 7088
S3.5/2 symbioticTiO 7054
Richard Walker 2011/03©Telluric O2
TiO 6815
TiO 6715
TiO 6681
TiO 6651
TiO 6478 TiO 6569
ZrO 6475/81/94
TiO 6358
TiO 6180
TiO 6159
Na I 5890/96
TiO 5847
TiO 5814
TiO 5668
TiO 5448/51
TiO 5167
TiO 5003/20
TiO 4965
TiO 4761-63
TiO 4626
TiO 4584
Spectroscopic Atlas for Amateur Astronomers 103
23 Carbon Stars on the AGB
23.1 Overview and Spectral Characteristics
The final stage of the stellar evolution on the AGB is formed by the deeply red shining car-
bon stars, in most of the cases also Mira Variables or „Carbon Miras“. In the atmosphere of
a star, finally moving to the upper part of the AGB, the C/O ratio becomes ܥ⁄ܱ > 1. This
results in a carbon excess which accumulates in a circumstellar cloud, dominating now im-
pressively the star's spectrum. Thus, in the intermediate class SC, and increasingly in the
following C–Class, moderately high resolution spectra show now predominantly absorp-
tions of diatomic carbon molecules. In addition to CH and CN the so-called Swan bands due
to C2 are particularly striking – discovered in 1856 by the Scot William Swan. (see also Ta-
ble 110). Further visible are atomic lines of S–Process products and impressive absorptions
of Na I. Angelo Secchi was the first to discover that the intensity gradient of the C2 Swan
bands is reversed [ ] in contrast to other molecular absorptions, such as titanium-
and zirconium oxide. For this feature, he created the separate spectral class lV. (see appen-
dix 34.3).
23.2 Competing Classification Systems
The phenomenon of carbon stars is still far from being fully understood (see e.g. [106]). For
the S–class, in spite of ongoing disputes it exist a generally accepted and consistently ap-
plied classification system. However for the carbon stars – the situation is still confusing
and unsatisfactory. The “Revised MK-System 1993”, propagated in most of the textbooks,
is applied obviously rarely. Its precursor from the 1960's, the so-called MK–C system, how-
ever, very often! Surprisingly frequent one can see even classifications according to the
much older Harvard system, with the R and N classes.
23.3 The Morgan Keenan (MK) – C System
This simple, old classification system is still very popular, both in many professional publi-
cations, as well as in most of the stellar databases. It uses the following format:
CX, n
X: defines on a scale of 0 – 7, the position of the star in the temperature sequence.
This scale is temperature equivalent to the spectral classes from G4 to M4 (see
table below). The system appears to have been extended down to C9 (e.g. WZ Cas).
n: This index rates on a scale from 1 – 5, the intensity of the C2 Swan bands. In
individual cases, appropriate supplements may be added, for example (e) for
emission lines, further also intensive lines of S–Process elements.
C0 C1 C2 C3 C4 C5 C6 C7
G4-G6 G7-G8 G9-K0 K1-K2 K3-K4 K5-M0 M1-M2 M3-M4
4500 4300 4100 3900 3650 3450
This classification system squeezes the entire complex class of carbon stars in one seven-
stage temperature sequence, supplemented just with one single intensity index!
Example: C 2,5 indicates a stellar temperature, equivalent to the spectral G9 – K0, com-
bined with very intense C2 Swan bands.
Spectroscopic Atlas for Amateur Astronomers 104
23.4 The „Revised MK System 1993“
In 1993, the old "C system" was revised to the "Revised MK System 1993" and adapted to
new findings [107]. It comprises five sub-classes (!), ordered by spectral symptoms –
whose astrophysical background however remains widely unclear. Perhaps this could be
one reason why the acceptance of this complex system, even some 20 years after its intro-
duction, seems to be still limited. The "Keenan 1993 system” uses the following format:
C–Sub X n
Sub: corresponds to the subclass of the carbon star according to the following table
X: defines the position of the star in the temperature sequence of the C–Class. This
forms a parallel sequence to the spectral classes from G4 to M4 (see table)
n: Various indices as listed below.
Sub- Supposed
class Criteria, spectral symptoms
status [106]
C–R Intrinsic Intensive Swan bands (C2) in the blue part of the spectrum. Considerable
flux in the blue/violet area – decreasing with lower temperatures. Later
C–N Intrinsic types show weaker H – Balmer lines, Hβ line serves as a temperature
indicator. S–process elements show average intensity. Intensity of
C–J Intrinsic 12C13C Band head at λ 4737 is above average.
C–H Extrinsic
S–Process elements are exceptionally intensive. Early C–N types show a
C–Hd ? tendency to Merrill Sanford Bands (SiC2 silicon carbide).Strong, broad
and diffuse absorption, or even a barely detectable flux in the blue range
λ < 4400. Swan bands C2 weaker than with the C–R types
Very intensive Swan bands (C2) and CN absorptions, further Merrill San-
ford bands (SiC2) and an infrared excess.
Dominant CH absorptions in the blue/violet area. Fraunhofer G band at
λ 4300 exceptionally- and S–Process elements above average intensive.
Hd indicates Hydrogen deficient. Hydrogen lines or CH absorptions weak
or even absent. CN and Swan bands (C2) above average intensive, very
often irregular R Crb Variables, showing sudden dramatic brightness
dips, due to veiling of the photosphere by circumstellar dust clouds.
Temperature equivalent ~Temperature C–R- C–N- C–H-
of the Main Sequence range [K] Sequence Sequence Sequence
4500 C–R0 C–H0
G4 – G6 4300 C–R1 C–N1 C–H1
G7 – G8 4100 C–R2 C–N2 C–H2
G9 – K0 3900 C–R3 C–N3 C–H3
K1 – K2 3650 C–R4 C–N4 C–H4
K3 – K4 3450 C–R5 C–N5 C–H5
K5 – M0 C–R6 C–N6 C–H6
M1 – M2 ~2500 C–N7
M3 – M4 C–N8
M5 – M6 C–N9
M7 – M8
Spectroscopic Atlas for Amateur Astronomers 105
This classification can be supplemented with the following indices [107], [2]:
Index Specification
C2 – Index Intensity of the molecular C2 Swan bands Scale 1 – 5
CH – Index Intensity of the molecular CH absorption. Scale 1 – 6.
MS – Index Intensity of the Merrill Sanford Bands (SiC2), Scale 1 – 5
J – Index
Intensity ratio of the C2 molecular absorption with the Isotopes 12C13C and
12C12C, Scale 1 – 5.
Elements In some cases, for strong lithium and sodium lines index values are speci-
fied
Not included in this system is the still not understood class of the dC carbon dwarf stars
located on the Main Sequence.
23.5 Function of the Subclasses in the Evolution of Carbon Stars
Current professional publications provide a rather diffuse picture about the functions of
these subclasses within the evolution of carbon stars. Many details, so e.g. about the
"dredge up" processes, are obviously far from being fully understood. Further the theoreti-
cal modelling of "carbon atmospheres” ܥ⁄ܱ > 1 seems to be difficult.
Analog to the S–Class also carbon stars on the AGB are referred as "intrinsic". It’s assumed,
that the subclass C–N, as well as the late representatives of C–R, form the last develop-
ment stage of the AGB sequence. These subgroups show a behaviour like M-giants with
very similar spectra. Furthermore they are also so-called "Low Mass Objects" with ≤ 3 solar
masses, typically show a mass loss and all of them are Variables.
Early representatives of C–R Type shows in contrast, neither variability nor a significant
mass loss. Further they seem not to be AGB stars and show a behaviour like K-giants. One
of the discussed scenarios postulates a star on the Horizontal Branch (HB) with a helium
“burning” core. However it remains here a mystery how the formed carbon will be conveyed
to the surface. A further hypothesis is a former binary system whose masses fused into one
single star.
The position of the C–J class on the AGB is unclear. It’s argued that these "Low Mass Ob-
jects" could be descendants of the C–R class on the AGB.
The C–H class seems to be "extrinsic". Most representatives are on the Horizontal Branch
(HB) and components of close binary stars. This suggests a mass transfer scenario similar
to the extrinsic S–Class.
23.6 Merrill Sanford Bands (MS)
Already at the beginning of the 20th Century in spectra of certain carbon stars, in the range
of some λλ 4640 to 5200, intensive absorptions attracted attention, which for a long time
could not be interpreted. The most intense band heads are at λλ 4977, 4909, 4867, 4640
and λ 4581. These absorption bands are named after the two explorers, who first described
these in 1926. Not until 1956 Kleman could prove with laboratory spectra that these bands
are caused by SiC2 silicon carbide. For this purpose he heated silicon in a graphite tube up
to 2500 K.
P. J. Sarre et al. have also shown [109] that the Merrill Sanford Bands are generated in
cooler layers of the stellar atmosphere, far beyond the photosphere. Merrill Sanford Bands
are most common in the entire C–J and the early C–N classes.
Spectroscopic Atlas for Amateur Astronomers 106
23.7 Commented Spectra
The majority of carbon stars is part of the sub-class C–N, documented here with some ex-
amples. The spectral classes are here indicated in the old MK–C system, according to Ste-
phenson's 1989 catalog [500]. If available [107], also the classification, corresponding to
the “Revised MK System 1993" is specified.
Table 64: A montage of three broadband spectra (200L grating) is presented to demon-
strate the difference between the profiles of differently classified carbon stars.
WZ Cassiopeiae: HD 224855, spectral class C9,2 Li or C-N7 III: C2 2 Li 10.
J2000 RA: 00h 01‘ 16‘‘ Dec: +60° 21‘ 19“ Vvar ≈ +6.5 – 8.5m
WZ Cas (1,500 ly) is an extremely late classified, cool Supergiant
(2,500 K). It’s the dominant component of a binary star system, im-
pressively contrasting with its white-bluish shining companion star
(mV ≈ 8m, A0). It’s located in the constellation Cassiopeia (image:
www.astro.sci.muni.cz/).
Besides the rather weak Swan bands (here classified with index value
2), this spectrum is dominated by the striking, almost fully saturated
(!) Na I line. Further by the impressive absorption of lithium Li I (λ 6708), whose intensity is
rated here with 10. Therefore WZ Cassiopeiae is often called "Lithium star" [348]. Accord-
ing to [349] this intense Li l absorption line was the first evidence of Lithium outside the
solar system, found by McKellar 1941 – a small but anyway interesting detail in the history
of science! Further the profile is dominated by CN- und C2-absorption bands and the H-
Balmer lines are barely recognisable here. The spectrum was recorded with the Celestron
C8 – exposure: 3x85 sec. The line identification is based amongst others on [100], [104],
[110], [348], [349].
Z Piscium: HD 7561, spectral class C7,3 or C-N5 C2 4
J2000 RA: 01h 16‘ 05‘‘ Dec: +25° 46‘ 10“ Vvar ≈ max. +6.8m
This star is a Supergiant with a surface temperature of some 3,000 K and located some
1,500 ly distant in the constellation Pisces [105]. With C7,3 it’s classified earlier than WZ
Cas (C9.2). The carbon absorptions are here much more intense and the Na l line some-
what weaker, but still very impressive! Instead of the exceptional Lithium line (λ 6708) the
undisturbed CN absorption at λ 6656 is visible here. The spectrum was recorded with DA-
DOS 50μm slit and the Celestron C8/ 3x60sec. The line identification is based amongst
others on [100], [104], [110], [348].
W Orionis: HD 32736, spectral class C5,4 Vvar ≈ max.+5.88m
J2000 RA: 5h 05‘ 59“ Dec: +1° 11‘ 27“
The luminosity class of this carbon giant (some 700 ly distant), is difficult to determine
[506]. Merrill Sanford Bands are striking here. In the profiles of the two other, much later
classified carbon stars, these absorption bands of triatomic SiC2 silicon carbide are hardly
recognisable. The spectrum was recorded with DADOS 25μm slit and Celestron C8/ 5x42
sec. The line identification is based amongst others on [100], [104], [110], [107], [109]
[348].
Table 64 A: W Orionis. Higher resolved spectrum (900L grating) in the wavelength domain
of the Merrill Sanford Bands. The line identification is based here additionally on [108].
Table 64 B: R Leporis, Hind's Crimson Star, HD 31996 spectral class C 7.6 e,
J2000 RA: 4h 59' 36" Dec: -14° 48' 23" mV = var. max +5.5m
R Leporis is probably the most famous representative of the carbon stars, discovered in
1845 by John Russell Hind. With a distance of 1100 ly, and a temperature of about 2290,
it is located in the constellation Hare. It is almost equally classified as Z Piscium, except of
the index e, which documents that the Hα line appears in emission here.
Spectroscopic Atlas for Amateur Astronomers 107
TABLE 64
CN 7283
CN 7119
CN 6955
WZ Cassiopeiae
Li l 6708 C 9.1
CN 6656 Z Piscium
Ca 6572 C 7.3
CN 6502 W Orionis
C 5.4
CN 6355
WZ Cas
CN 6259 Z Psc
CN 6206 W Ori
C2 6168
C2 6122 Richard Walker 2011/04©
C2 6059
C2 6005 C2 6191
Na l 5895.92
Na l 5889.95
CN 5749
C2 5636
C2 5585
C2 5541
C2 5502
Ti l 5426.26
Mn l 5394.67
Sc l 5301.94
CN 5255
Cr l 5206.04
C2 5165
C2 5130
Ba ll 4934.1 SSiiCC22Merill Sanford4977
Bands 4957
I=0.0SiC2 4909
SSSiiiCCC222 4867
4832
4807
C2 4737 SiC2 4767
Spectroscopic Atlas for Amateur Astronomers 108
TABLE 64A Richard Walker 2011/04©
C2 /Sc l 5302
CN 5255
Fe I/Ni I/V II 5215-16
Fe I/Cr l/La II 5205
SiC2 5198
SiC2 5192
C2 5165
C2 5130
C2 5095-110
W Orionis C 5.4 Merrill Sanford Bands
C2 5041 SiC2 4981
Merrill Sanford Bands C2 5020
C2 5006
SiC2 4977
SiC2 4957
Ba ll 4934.1
SiC2 4909
SiC2 4867 I=0.0
SiC2 4851
SiC2 4832
SiC2 4807
SiC2 4767
SiC2 4750
C2 4737
Spectroscopic Atlas for Amateur Astronomers 109
TABLE 64B CN 7283 Richard Walker 2014/02©
CN 7119
CN 6955
CN 6656 Hα 6562.82
Ca 6572
CN 6502
R Leporis „Hind‘s Crimson Star“ HD 31996 C 7.6 e CN 6355
CN 6259
CN 6206
CC22 6191
6168
C2 6122
C2 6059
C2 6005
Na l 5895.92
Na l 5889.95
CN 5749 I=0.0
C2 5636
C2 5585
C2 5541
C2 5502
Ti l 5426.26
Mn l 5394.67
Sc l 5301.94
CN 5255
Cr l 5206.04
C2 5165
C2 5130
Spectroscopic Atlas for Amateur Astronomers 110
24 Post AGB Stars and White Dwarfs
24.1 Position of Post AGB Stars in the Stellar Evolution
This section explains the final stage for stars with less than about 8 solar masses. This is
too light for a final SN explosion of type II. Only as a member of a binary system a White
Dwarf can become a SN type Ia by accretion of matter and finally exceeding its critical
Chandrasekhar mass limit (sect. 25).
24.2 Post AGB Stars
After the final AGB-stage as a carbon star it begins to repel its envelope as a planetary neb-
ula. This stellar stage is called "post-AGB" and includes also the central stars of planetary
nebulae. Such a very early Post-AGB-Object (HD44179) is presented in sect. 28 as an excit-
ing source of the protoplanetary nebula Red Rectangle (Table 85).
24.3 Spectral Features at Post AGB Stars
During the repulsion of its shell the star, now increasingly becoming hotter, performs an
impressive loop in the upper part of the HRD and passes on this "farewell tour" almost all
spectral classes (see chart sect. 20). In extreme cases, its temperature may reach far be-
yond 100'000K and the dying star can generate for a very short time even a Wolf Rayet-like
spectrum WRPN (sect. 28.3).
24.4 White Dwarfs
By pushing off of a planetary nebula, the star loses mass, the thermonuclear fusion proc-
esses inside the star extinguish and the remaining rest is finally reduced to an earth-sized,
extremely dense object, with an enormously strong gravitational acceleration at its surface.
Most of the White Dwarfs are composed of a "degenerate" carbon-oxygen core, the prod-
ucts of the previous helium fusion, and a thin shell of hydrogen and helium. The spectral
class sinks in the HRD to its lowest area of the White Dwarfs. The absorption spectrum is
produced here just by the remaining residual heat of the very slowly cooling stellar corpse.
Their absolute luminosity is now so low, that for amateurs only a few objects in the imme-
diate solar neighbourhood are reachable, ie within a radius of about 50 light years. A corre-
sponding list can be found in [262]. The nearest and brightest With Dwarfs are the com-
panion stars Sirius B and Procyon B. Anyway, due to their close orbits around the A-
components they are spectroscopically inaccessible for amateurs. Easiest to observe, visu-
ally and somewhat limited also spectroscopically, is 40 Eridani B (mV = 9.5) as a component
of a triple system, identified as a white dwarf not until 1910. Visually, this object was de-
tected by William Herschel already 1783. The brightness of the remaining White Dwarfs
lies already within the range of magnitudes 12m – 13m.
24.5 Spectral Characteristics and Special Features of White Dwarfs
The exorbitantly high gravitational acceleration at the surface causes, especially due to the
Stark effect of the interatomic electric fields (Kuiper 1939), extremely broadened absorp-
tion lines. This effect forms here the spectral "brand" and affects mainly the hydrogen
Balmer series. In addition, also lines of helium and calcium may appear. Due to the low
brightness the display of the finer absorption lines remains here reserved to high-resolution
spectrographs at large professional telescopes.
In astrophysics, the gravitational acceleration g is expressed as a logarithm to the base 10,
however strangely not in ሾ݉/ݏଶ], but in [ܿ݉/ݏଶ]. For Sirius B this value is log g = 8.57 , cor-
responding to some 371ᇱ000ᇱ000 ܿ݉/ݏଶ or 3′710′000 ݉/ݏଶ. Compared to just 9.81 ݉/ݏଶ on
earth, this is so extremely high, that the gravitational redshift, predicted by Einstein's Gen-
Spectroscopic Atlas for Amateur Astronomers 111
eral Theory of Relativity, becomes relevant even for amateurs [30]. The required work for
the light (photons) to escape the gravitational field of a White Dwarf, causes at Hα a red
shift of slightly <1Å, corresponding to a radial velocity of about 20–40 km/s. This star-
specifically different value must therefore inevitably be accounted for measurements of the
Doppler shift. The log g values of white dwarfs are in the range of approximately 7–9.
24.6 Classification System by McCook & Sion
The first letter of the classification is D, which means "Degenerate". The second letter indi-
cates the primary spectroscopic characteristics [2], [261]. Besides the type DA the physical
background of most other subclasses is not yet fully understood.
DA – Only Balmer lines; no He I or metals
DB – He I lines; no H or metals present
DC – Continuous spectrum, no lines deeper than 5% in any part of the spectrum
DO – He II strong; He I or H present
DZ – Metal lines only; no H- or He lines
DQ – Carbon features, either atomic or molecular in any part of the spectrum
Additional letters and symbols may specify further effects and features:
P – Magnetic white dwarfs with detectable polarisation
H – Magnetic white dwarfs without detectable polarisation
X – Peculiar or unclassifiable spectrum
E – Emission lines are present
? – Uncertain assigned classification
V – Optional symbol to denote variability
d – Circumstellar dust
C I, C II, O I, O II – atomic species in spectra of hot DQ-Dwarfs
The classification digit, following this letter combination, specifies the effective tempera-
ture ܶ according to the formula:
ܶ = 50ᇱ400/ݏݏ݈ܽܥ. ݀݅݃݅ݐ
For rarely occurring temperatures above 50'400K the classification digit becomes less than
1.0 and is expressed for example as .9 .8 .7 etc.
Finally sometimes also the log g value of the gravitational acceleration is added. So the
classification eg DA 2.5_7.8 means a spectrum with hydrogen lines, ܶ of about 20'000K
and a gravitational acceleration at the surface of log ݃ = 7.8.
24.7 Commented Spectra
Table 65: shows a montage with broadband spectral profiles (200L/mm) of three differ-
ently classified White Dwarfs, all normalised on the same continuum section. Depending on
the source consulted, here [260], the indicated spectral class may vary considerably.
WD 0644 +375: GJ 246 Spectral class DA 2.3 Distance 50 ly [262]
J2000 RA: 06h 47' 38" Dec: +37° 30' 57" mV = 12.1m
Besides the highly broadened hydrogen lines no further spectral features are visible here.
The equivalent width of the Hβ line is here with 65Å about 2.4x as large as by the already
impressive absorptions of the main sequence star A1Vm Sirius (27 Å). The classification
digit (2.3) suggests here an effective temperature of about 22'000K. The spectrum was re-
corded with the 50μm slit and the 200L grating. Exposure time: C8/DADOS/Atik 314L+:
1x1800 sec, 2x2 Binning Mode, –20°C.
Spectroscopic Atlas for Amateur Astronomers 112
WD 0413 –077: 40 Eridani B Spectral class DAP 3.1 Distance 16 ly
J2000 RA: 04h 15' 22" Dec: –07° 39' 29" mV = 9.5m
Here again, the strongly broadened hydrogen lines are the most prominent spectral feature.
The equivalent width of the Hβ line is here with 78Å even about 3x as large as by the any-
way impressive absorptions of the main sequence star A1Vm Sirius (27 Å). Compared to
WD 0644 +375 this is caused here by the lower temperature of 40 Eridani B, which fa-
vours this type of absorption and may be estimated, in accordance with the classification
digit, to about 16'000K. The following figure shows the superimposed profiles of 40 Eridani
B (red) and Sirius (blue). It is striking that the broadening chiefly affects the Hβ-and to a
lesser extent also the Hγ line.
On Table 65 the Hα line at 40 Eridani B appears strongly deformed. Whether this effect is
caused by the strong magnetic fields (additional letter P in the classification) is not clear.
The figure on the right shows the B-component (mV = 9.5m) of A Spalt
the triple star positioned at the lower end of the 50 μm slit.
B
Thereby it was avoided that the spectrum has been contami- C
nated by the nearby C-component (mV = 11.2m). The A-
component (mV = 4.4m) is here the Guide Star at a comfortable
distance of 83". Recording: DADOS 50μm slit, Celestron C8/
2x1800 sec.
WD 0046 +051: Van Maanen 2, Van Maanen‘s Star Spectral class DZ 8
J2000 RA: 00h 49' 10" Dec: +05° 23' 19" mV = 12.4 Distance 14 ly
Adriaan Van Maanen discovered this object in 1917 as the first stand alone White Dwarf.
This object has already been cooled down to about 6000K, a similar temperature range as
the solar photosphere. Its surface is possibly strongly contaminated with interstellar and
planetary particles (metals) [265]. Therefore at this resolution, just the two intense and
strongly broadened Fraunhofer H and K lines of ionised calcium Ca II can be seen here. Re-
cording: DADOS 50μm slit, Celestron C8/ 1x1800 sec.
White Dwarfs Spectroscopic Atlas for Amateur Astronomers
Hβ 4861.33 WD 0644 +375 DA 2.5 TABLE 65
Hγ 4340.47 WD 0413 -077 40 Eridani B DA 4 Hα 6562.82
Hδ 4101.74
Hε 3970.07 WD 0046 +051 Van Maanen Star DZ 7
Ca II 3968.47 WD 0644 +375
Ca II 3933.66 WD 0413 +077
WD 0046 +051
113
©Richard Walker 2014/02
Spectroscopic Atlas for Amateur Astronomers 114
25 Supernovae
25.1 Phenomenon of Supernova Explosion SN
A supernova explosion totally destroys the star and
forms the definitive end point in its life. By this
cataclysmic runaway reaction an unimaginable
amount of energy is set free and almost the entire
stellar mass, initially with> 10,000 km/s, is distrib-
uted to the surrounding space. For comparison: the
detonation velocity of our most rapid explosives just
reaches ~ 8km / s (Nitropenta). As a result of such
an explosion, the interstellar matter (ISM) is en-
riched with heavy elements, which decisively influ-
ences the later formation of stars, planets and fi-
nally also of possible life. The diameter of old Supernova Remnants (SNR) may finally reach
up to some 100 ly, so eg the famous Cygnus Loop. Otherwise the diameter of the relatively
young Crab Nebula M1 is just about 11 ly (see sect. 28). The image (HST) shows the SN
1987A (SN type II) in the Large Magellanic Cloud (distance ~168'000 ly) – about 20 years
after the apparent explosion time.
25.2 Labelling of Supernovae
Supernovae are labelled with the letters SN, followed by the year of discovery and an ongo-
ing assigned letter, such as SN 2014 J. Since several hundred SN are discovered each year
with today's telescopes and automatic monitoring systems, after the first 26 events, the
assigning of double letters becomes necessary.
25.3 Classification of SN Types
SN are divided into the two main types, labelled with the roman numerals I and II. This
rough subdivision is very easy even for amateurs, since the spectra of SN type II show
emissions of the H-Balmer series – and those of SN type I show none. This quite simple re-
lation was discovered as early as 1941 by Rudolph Minkowski.
The characteristic lack of hydrogen in SN type I is caused by two very different scenarios,
and therefore the division into the sub-classes Ia and Ib / Ic is required:
Type Ia: For stars with ~<8 M the hydrogen envelope is repelled as a Planetary Nebula.
So here all former main-sequence stars of spectral classes later than about B4-
B6 are concerned. What finally remains is a White Dwarf (sect. 24), which in
most of the cases consists mainly of carbon and oxygen.
Type Ib/Ic: For stars with ~>25 M [234] the hydrogen envelope is repelled as a Wolf
Rayet Nebula. So here roughly all former main-sequence stars of spectral class
O are concerned. What finally remains is an extremely hot Wolf Rayet Star (sect.
9), at which SN type Ib shows helium lines and SN type Ic shows none.
For the SN type II, with the characteristic hydrogen lines in the spectrum, it remains just the
huge middle mass range of ~8–25 M . Otherwise considering the spectral class this area
is rather small and concerns roughly all former main-sequence stars in the rough area of
just the early B- to the late O-class! A significant contribution to this theory stems from the
Swiss astronomer Fritz Zwicky in the 1960s.
SN show also a few outliers, which are not yet fully understood. Thus, e.g. in rare cases an
SN may start as Type II but end up as Type Ic [2]. Generally speaking for SN, by far not yet
all relationships are fully understood here.
Spectroscopic Atlas for Amateur Astronomers 115
25.4 Explosion Scenario "Core Collapse"
The core-collapse-SN forms the end of all stars with >8 M . This theory was proposed by
Fritz Zwicky in 1938. At the end of the giant stage, increasingly heavier elements are gen-
erated in the core of the star. Finally it will be fatal for the star, when it starts to produce
iron. For the formation of this, and all the following even heavier elements, the fusion proc-
esses consume energy. Thus, due to lack of radiation pressure, the core can no longer re-
sist the gravitational forces and therefore it inevitably collapses. This initial implosion is
then quickly changing in to a cataclysmic explosion and at the very end it remains an ex-
tremely dense object with just a few km in diameter. Depending on the original stellar mass
either a neutron star or a black hole is formed. During the core-collapse, the very rare, ex-
tremely large stars of the Wolf Rayet type can further eject a highly intense gamma ray,
headed parallel to the stellar rotation axis (Gamma Ray Burst). In extreme cases it may oc-
cur a hypothetical Hypernova such is expected for Eta Carinae. The physical effects of the
shock wave within the expanding SNR are described in sect. 28.
25.5 Explosion Scenario "Thermonuclear Carbon Fusion"
This scenario is also called thermonuclear explosion and is exclusively limited to the
SN type Ia with an initial stellar mass of <8 M . As a member of a binary system, a White
Dwarf (sect. 24) by accretion of matter and finally exceeding the critical Chandrasekhar-
mass limit (ca. 1.4 M ), may explode as SN type Ia. However as an additional condition a
minimum rate for the annual accretion must be exceeded (߂݉/ > ݎܽ݁ݕ1 ∙ 10ି଼ M ) [297],
otherwise mostly just recurrent Nova explosions occur. However, such much smaller events
are limited to the currently accreted material at the stellar surface, in which the star itself
will not be destroyed. But if all conditions are met, the degenerated electron gas can no
longer withstand the gravitational pressure and the stellar core, mostly consisting now of
carbon and oxygen (C/O), explodes. This cataclysmic event is caused by the sudden onset
of nuclear carbon-fusion, which is why SN Type Ia is sometimes also referred as Carbon
Detonation Supernova. Anyway in contrast to the Core-Collapse SN, SN Type Ia leaves no
residual object.
25.6 SN Type Ia – Standard Candle
As already mentioned the SN type Ia exclusively occurs with White Dwarfs if they exceed by
accretion the "quasi standardised" critical Chandrasekhar-mass limit of about 1.4 M . Thus,
this way a more or less uniform amount of energy of about 10ସହ J is released [295]. Further,
both the photometric and the spectral profiles of such events are very similar. With
SN Type Ia the luminosity reaches the maximum after about 20 – 30 days, with an absolute
magnitude in the blue region of MB ≈ –19.5M [296]. These values remain within just a small
stray area and are, for SN explosions, clearly in the region of the top rankings, which is
bright enough to outshine an entire galaxy for several months. Due to this uniform appear-
ance the SN Ia explosions serve also as indispensable "standard candles" for measuring the
entire observable universe.
For this purpose the core-collapse SN types are less useful, because the intensity of such
explosions depends strongly on the initial stellar mass. The absolute maximum brightness
lies here within a very large stray area of MB ≈ –15M to –21M, see [296, Fig. 2]. The ex-
treme values are denoted here as sub luminous and over luminous [296].
25.7 Spectral Determination-Diagram for the SN Type
The following diagram is used for the spectroscopic identification of the SN type. In addi-
tion, the relevant explosion scenario, the type of the progenitor star, and finally the rough
order of magnitude for the original stellar mass are assigned here.
Spectroscopic Atlas for Amateur Astronomers 116
Not shown here are the subcategories of type II, which are determined photometrically by
the course of the light curve (see [2]). The light curve of Type II-P shows a plateau phase
after the maximum while the brightness of the Type II-L ~ decreases rather linearly. Another
Subtype II n shows narrow lines.
Spectral Features Mass of Progenitor Explosion
the Star Before Scenario
No Hydro- Yes Original Explosion
gen ? Star
Yes Si II at No [Solar
6150 Å(1) ?
masses M☉]
Yes He I at No
5876 Å(1) ?
Ia M < 8M☉ White Dwarf Carbon
with 1.4 M☉ Fusion
SN II 8M☉< M <25M☉ Red Giant Core
Type Collapse
M >25M☉
Ib
Wolf Rayet
Star
Ic M >40M☉
(1) These values are subjected to the Doppler shift and may vary
25.8 SN Type Ia – Features in the Optical Spectral Range
The spectral features of SN type Ia are demonstrated here by the example of SN 2014 J. On
this topic exist, mainly by A. Filippenko and D. van Rossum, excellent and well readable
publications, even for advanced amateurs, (see bibliography sect. 40). With coloured
charts, based on model calculations, D. van Rossum impressively shows [291], which
highly complex blends form the SN type Ia-profiles. Both, the distinction between absorp-
tions and emissions, as well as the determination of the real continuum-course, is ex-
tremely difficult here. So the relative Flux-calibration of the spectrum, with a standard star,
makes really sense here [30].
Accordingly, "absorptions" are consequently called "troughs" here and their equivalent
widths EW denoted as "pseudo-EW" [pEW] [291]. The ions, labelled here on the profile, are
only the detected (and in some cases merely suspected) main causes of such "structures".
Another quotation (just roughly transferred): "Apparent absorptions are often just random
gaps between two emissions".
Spectroscopic Atlas for Amateur Astronomers 117
Table 67: SN 2014 J, Type Ia, Host Galaxy M82
SN 2014 J was since decades the brightest super-
nova and reached about January 30th, 2014, the
maximum apparent brightness of mV = +10.5. It
was discovered by chance and surprisingly late not
until January 21st, on the occasion of a student
exercise at the University of London. The apparent
explosion date was subsequently determined from
pictures, taken by automatic surveillance systems,
to January 14th 2014 [292].
The picture on the right was taken by Roland Stal-
der at the Hubelmatt Observatory in Lucerne [711]
and shows in the crosshairs the SN 2014 J within the host galaxy M82 [711]. Here, on
January 31st, 2014, 17 days pe (post explosion), also the upper, red profile on Table 67
was recorded. It shows the usual SN Type Ia spectrum near the maximum brightness with
the typical two sulphur sinks (S II) at λ 5400, also called "W-absorption", and the prominent
silicon trough ("Silicon Valley") at λ 6150. The latter forms the key feature for the identifica-
tion of the SN type Ia which allows also the rough estimation of the detonation velocity.
This huge Si II silicon absorption with the rest wavelengths of λλ 6347 and 6371 appears
here impressively blue-shifted by about 200 Å to the range of λ 6150. Evaluated with the
spectroscopic Doppler law [30] it results a radial velocity of ܿ ∙ ≈ ݖ9ᇱ800݇݉/ݏ. A similar
value results also with the Doppler law [30], applied to the FWHM of the Si II-trough at λ
6150. A. Filippenko and D. Van Rossum [291] call this the "Si II-velocity".
Approx. 20 days later, ie 37 days pe, in the blue bottom profile, as expected, several Fe II
emissions are prevailing and suppressing for example the S II "W-absorption" at λ 5400.
However the impressive "silicon trough" is still present here showing still a comparable in-
tensity. A striking feature of both profiles is the intense sodium Na I line, which is corre-
spondingly interpreted by several publications as interstellar absorption within the appar-
ently extremely dusty host galaxy M82. A. Filippenko analyses this line even as saturated
[292]!
The two relatively flux calibrated profiles [30] show a slight intensity-increase towards the
long-wave (red) direction. This is another strong evidence for the huge interstellar redden-
ing because the "unreddened" model spectra [291] show the intensity peak in the UV
range.
Recording info Red Profile: January 31st, 2014: Hubelmatt Observatory Lucerne [711],
40 cm MFT Cassegrain, focal length 5550mm, DADOS with 200L/mm, 50μm slit width,
Atik 314L+ 2x900s, 2x2 Binning.
Recording info Blue Profile: February 20 2014: Celestron C8, DADOS with 200L/mm,
50μm slit width, Atik 314L+ 1x1800 sec, 2x2 Binning.
SN 2014 J M82 Estimated first light SN 2014J: Jan. 14.75 UT ~JD 2456672.3 Spectroscopic Atlas for Amateur Astronomers
(Filippenko et al.) pe = [days] „post explosion“
Spectrograph: DADOS, Grating 200L/mm, Slit: 50µm TABLE 67
Profiles relatively Flux calibrated 31.1.2014 , 2030 UTC , JD 2456689.35 +17 pe
20.2.2014 , 2030 UTC , JD 2456709.31 +37 pe Telluric O2 6870
Na I
? Fe II Fe II S II
Ca II Fe III Fe III
Si II Si II
?
Mg II Fe II Si II Fe II
Fe II
©Richard Walker 2014/03
Fe II 118
Spectroscopic Atlas for Amateur Astronomers 119
26 Extragalactic Objects
26.1 Introduction
It’s impossible with amateur equipment, to record spectra of single stars within external
galaxies. However it’s feasible to record composite spectra of galaxies and Quasars! In con-
trast to profiles of individual stars the composite (or integrated) spectra show the super-
posed characteristics of hundreds of billions individual star spectra. Using Doppler spec-
troscopy thus also the radial velocities, respectively the z-values of such objects can be
measured. Further by "edge on" galaxies the rough distribution of the rotation speed within
the galactic disk can be estimated [30].
On a professional level, this has been practiced successfully since the beginning of the
20th century and has contributed substantially to our present understanding of the uni-
verse. The first, who tried this with M31, was Vesto M. Slipher in 1912 at the Lowell Ob-
servatory in Flagstaff Arizona. He was able to measure the blue shift of the spectrum and
derived a radial velocity of –300 km/s. Further he detected the rotation of this "nebula".
The fact that M31 is a galaxy outside the Milky Way was proved only later in the early
twenties. Further, he also noted that most of the other galaxies appear red shifted and thus
are removing from us. Lemaître and Hubble used these shift measurements later for the
correlation with the distance (Hubble constant).
26.2 Morphological Classification
The characteristics of the spectra are partly correlated with the morphological types of gal-
axies. The following graph shows the so-called Hubble Sandage tuning-fork diagram.
Image: NASA/ESA
It is based on the former faulty hypothesis that this sequence should represent the evolu-
tion of galaxies, starting from the elliptical shape of E0 and ending with the spiral types Sc,
or SBc. Similar to the stellar spectral classes, we therefore call, unfortunately even today,
the elliptical as "early" and the spirals as "late" types. However, according to current knowl-
edge, the elliptical galaxies represent rather the final stage, merged by a number of smaller
galaxies. During this process also the irregular stage Irr is passed through. Examples are
M82 and the Antenna Galaxy NGC4038/4039.
Spectroscopic Atlas for Amateur Astronomers 120
26.3 Spectroscopic Classification
The spectroscopic classification distinguishes these objects according to spectral features.
Already in the 1920ies it was recognised that such composite- or integrated spectra have a
similar information content like those of individual single stars. They also show profiles e.g.
with or without emission lines, chiefly depending on the kind and activity of the galactic nu-
clei. Such spectral profiles are powerful means to determine e.g. the content of stars, and
the development state of the galaxy [301].The relative similarity of the galactic composite
spectra to those of individual single stars, was also a convincing argument in the historical
dispute (mainly between H. Shapley and H. Curtis), that galaxies are not dust or gaseous
nebulae like M42 (see Tables 80/80A), but extremely distant, huge clusters of individual
stars!
In the second half of the 20th Century, classification systems have been developed, which
were based on the similarity with stellar spectral types. The so called Integrated Spectral
types of the galaxies were determined this way by W. Morgan, beginning with A, AF, F- and
ending by the late K-systems. The following correlations have been noticed [2]:
– The later the stellar spectral type, with which the profile shows similarities, the stronger
is the galaxy centrally concentrated.
– The elliptical classes E1 – E4, remain still dominated by absorption features of later
spectral types. From here on however, the characteristic changes until the profiles of
the types Sc/SBc and Irr look similar to those of early spectral classes. Furthermore
emission lines show up here with increasing intensity.
The following table shows another, also rather rough classification system, based on the
presence, intensity and shape of emission lines. It also includes peculiar shapes such as
Seyfert Galaxies and Quasars, standing out by an extremely high core activity and therefore
belonging to the category of AGN (Active Galactic Nuclei). The core activity increases here
from top to down.
Spectroscopic Atlas for Amateur Astronomers 121
26.4 Rough Scheme for Spectroscopic Classification
Galaxy Type Features Within the Affected Objects and Examples
Optical Spectral Range Responsible Effects
Absorption Almost exclusively absorption Mainly elliptical galaxies M31, M33,
line Galaxy lines, sporadically weak emis- with a rather old star in- M81
sions may show up ventory, low star forma-
tion rate and weakly ac-
tive nuclei
By the majority absorption Mainly types E0 – SO M94, M104
lines. In the core region emis- with a younger star inven-
LINER Galaxy sion lines of weakly ionised or tory compared with ab-
neutral atoms as O I, N I and S sorption line galaxies.
I. Possibly weak emissions of Slightly increased core
ions with higher ionisation en- activity.
ergy as He I, Ne II and O II.
Starburst Intensive H-emission lines, fur- Colliding, gravitationally M82,
Galaxy ther just weakly may appear: [N interacting, or gas-rich NGC
II] (λ6583), [S II] (λλ 6716/31) young galaxies with giant 4038/4039
and [O I] (λ6300), generally lit- H II regions and a corre-
tle or no absorptions spondingly high rate of
star birth
Seyfert Intensive and Doppler- Central, supermassive M77,
Galaxy broadened H- and slim [O III] black hole with high M106, NGC
emission lines, [N II], [S II], and accretion rate 4151
[O I] are intensive, generally
little or no absorptions
Intensive and extremely Dop- Central, supermassive 3C273
pler-broadened H- emission
Quasar lines black hole with extremely
high accretion rate
Continuum spectrum of the Blazars differ from qua- Makarian
bright matter jet without any sars just by our perspec- 421
spectral lines and irregular, tive; here we look directly
Blazar, strong intensity fluctuations. into the synchrotron ra-
BL Lacertae Spectral lines measurable only diation of the jet matter.
Object in the very faint ambient galaxy
disk with large telescopes.
Spectroscopic Atlas for Amateur Astronomers 122
26.5 Absorption Line Galaxies
The principle of the absorption line galaxy is demonstrated here by the example of the An-
dromeda galaxy M31. In this type, star-like spectra appear almost with exclusively absorp-
tion lines.
Table 70: Absorption line Galaxy Andromeda M31/NGC224
The table shows heliocentric parameters according to NED [501], and Evans, N.W. Wilkin-
son (2000) [1] (Value including suspected dark matter).
Radial Redshift z Distance Morphology Diameter Mass 1
Velocity
–300km/s –0.001 2.6 M ly Sb 141,000 ly 1.23 1012 M
Table 70 shows a composite spectrum of M31. The slit Slit
axis of the DADOS spectrograph was adjusted nearly par- axis
allel to the apparently smaller cross axis of the galaxy,
since no rotation effects are intended to show here [30].
According to [301] composite spectra of the galaxy types
Sa – Sb are dominated mainly by developed Giant stars.
As a matter of fact the pseudo-continuum of the recorded
M31 spectrum fits here best to that of a single star of the
late G–class what is called Integrated Spectral Type. For a comparison, the M31 profile is
therefore shown here superposed with the pseudo-continuum of Vindemiatrix (ε Vir).
M31 belongs obviously to the category of absorption line galaxies. The most dominant
common features are the Fraunhofer G-band, the Calcium Triplet, and the Sodium Double
Line (Fraunhofer D line). The H-Balmer lines are just very faintly visible in absorption. Ac-
cording to [301], this finding fits rather to elliptical galaxies (type E). For the type Sa–Sb
also few emission lines, caused by a younger stellar population, may sporadically be recog-
nisable. However this is not detectable in this spectrum of M31.
The profile was recorded in the Nasmith focus of the 90 cm CEDES Cassegrain Telescope in
Falera – exposure: 3x340 sec. Clearly visible here is the expected blue shift of several Ang-
stroms by the absolutely wavelength-calibrated M31 profile, compared to the relatively
calibrated spectrum of ε Vir, based on known lines.
26.6 LINER Galaxies
The very numerous so-called LINER galaxies form a kind of transition between absorption
lines and starburst galaxies. LINER means "Low-Ionisation Nuclear Emission-Line Region".
The spectral features of the galactic core region are emission lines of weakly ionised or
neutral atoms such as O I, N I and S I. Just weakly possibly appear emissions of highly ion-
ised atoms such as He I, Ne II and O II. In contrast to the Seyfert galaxies (sect. 26.9), the
core at LINER galaxies is still relatively faint. Well-known examples are M104 (Sombrero)
and M94. Morphologically mainly the types E0 – SO are concerned. This phenomenon, es-
pecially the origin of the involved ionisation processes is still subject of debate.
Spectroscopic Atlas for Amateur Astronomers 123
Table 71: LINER Galaxy M94
Heliocentric parameters according to NED [501] and [1]: Trujillo et al. [315].
Radial- Redshift ݖDistance Morphology Diameter Mass [1]
Velocity ݒ
+308km/s +0.0010 16.4 Mio ly SA 67‘000 ly 6.5 1010 M
This typical "Face On" galaxy (Image: HST) has achieved
general awareness because its rotational behaviour can
apparently be explained without the currently hotly de-
bated "dark matter". Both profiles of Table 71 have been
recorded in the integrated light at a time interval of 48
hours and (inevitably) with different slit positions relative
to the Galaxy. The upper, red profile may represent rather
the peripheral regions. Compared to M31 amazingly
many absorption lines can be identified here with almost
star-like shape. This composite spectrum shows strong
similarities with a stellar profile of the late F-class. Par-
ticularly striking here is one of their "brands" – the comparable intensity of the G-band (CH
molecular), and the directly adjacent Hγ line (sect. 16.3). This striking “Line-double” can ex-
clusively be seen in the F-class. This also fits to the intensity of the two Fraunhofer Ca II
lines and the still rather modest absorption of the Magnesium triplet (λλ 5167 -. 5183).
Compared to M31 the entire H-Balmer series appear here very prominently and also all
other symptoms indicate a much younger star inventory in M94. The lower, blue profile may
represent rather the weakly active core of the galaxy. Some possible, unidentifiable emis-
sion peaks are marked here with red arrows. The most striking features, as the unique dou-
ble peak emission at about λ 4507, are marked with red ellipses. For Hβ it remains unclear,
whether we see here double-peak absorption or a shell core emission. Recording info:
C8/DADOS/Atik 314L+: 1x1800sec, 2x2 binning mode.
26.7 Starburst Galaxies
Starburst Galaxies are mostly colliding, gravitationally interacting, or young, gas-rich galax-
ies with giant H II regions and a correspondingly high star birth rate. Clearly over-
represented here are the irregular types, such as M82 or the Antenna Galaxy NGC
4038/4039. The most striking spectral features are intense emissions of the H-Balmer se-
ries, the virtual absence of absorptions, just weak emissions of [N II] (λ6583), [S II] (λλ
6716/31), and possibly also [OI] (λ6300).
Table 72: Starburst Galaxy M82, Cigar Galaxy
Heliocentric parameters according to NED [501] and J.P. Greco, P. Martini et al. [1]
Radial- Redshift ݖDistance Morphology Diameter Mass [1]
Velocity ݒ
+203km/s +0.0007 12.6 Mio ly Irr 47‘000 ly ~1010 M
Spectroscopic Atlas for Amateur Astronomers 124
The apparent chaotic structure of this galaxy can be at-
tributed to gravitational interaction with the much larger
neighbouring galaxy M81, as well as effects of the spe-
cific perspective (Image: HST). M82 is a typical represen-
tative of the starburst galaxies with very few weak ab-
sorption lines. Na I is here most likely of interstellar ori-
gin, mainly within M82. Additional features include in-
tense emissions of the H-Balmer series and probably
shock wave induced sulphur lines. The forbidden [OIII]
lines are here barely recognisable. Also striking is here
the double peak in the range of the Hγ-emission (compare also Seyfert galaxy M77, Table
73). Compared with galactic emission nebulae (sect. 28) as well as M77, the excitation
level is very low here. The forbidden [OIII] line at λ 5007 shows here up just very weakly
and the Hα emission is much more intense compared to [N II] at λ 6583. Recording info:
C8/DADOS/Atik 314L+: 1x1800sec, 2x2 binning mode.
M82 is also the host galaxy of the brightest supernova since decades, SN2014j, type Ia
(see sect. 25, Table 67).
26.8 The phenomenon of AGN (Active Galactic Nuclei)
The AGN phenomenon in the subsequently described
Seyfert galaxies and Quasars is caused by so-called su-
permassive black holes, collecting vast amounts of mat-
ter from their surroundings and generating thereby exor-
bitant amounts of energy and intense X-ray radiation.
This process is accompanied by an accretion disk in the
equatorial plane and jets of matter which are ejected
parallel to the rotation axis of the object with almost the
speed of light (image: Wikipedia). Seyfert galaxies are the
largest group with AGN [307].
26.9 Seyfert Galaxies
The principle of the Seyfert galaxy is demonstrated here by the example of M77.
Table 73: Seyfert galaxy M77/NGC1068
The table shows heliocentric parameters according to NED [501] and other sources [1].
Radial Redshift ݖDistance Morphology Diameter Mass [1]
Velocity ݒ
+1'137km/s +0.0037 44 M ly Sb / Sy2 115‘000 ly ~1.0 1012 M
Table 73 shows a composite spectrum of the core of M77. The
galaxy is classified as type Seyfert 2 (image: NASA/Wikipedia).
In the 1940ies Carl Kennan Seyfert (1911 – 1960) discovered
in the core of some galaxies intensive emission lines of the H–
Balmer series with Doppler-broadenings of more than 1000
km/s.
In addition, emissions of forbidden transitions, such as [O III]
and [N II] can be detected. However, they can’t substantially be
Spectroscopic Atlas for Amateur Astronomers 125
broadened, due to the shock sensitivity of the metastable initial states. In contrast to the
emission of the H-Balmer series, they are probably generated far away from the turbulent
core around the supermassive black hole, whose mass is estimated to be ~15 Million M
[Hubble ESA, Garching]. In contrast to the permitted lines they show therefore virtually no
intensity fluctuations.
Also remarkable are the double peaks at the Hγ and the Ne III (3967) emission. B. Garcia-
Lorenzo et al. [304] and other authors suggest here Doppler effects due to differently run-
ning streams of gas in the close vicinity around the black hole.
Seyfert galaxies are divided into:
– Subclass 1 with strongly broadened lines, limited to the permitted transitions
– Subclass 2 with only slightly broadened lines, limited to the permitted transitions
Generally it is now assumed, that this difference in classification is rather caused by effects
of different perspectives, see "Seyfert Unification Theory" [306]. Thus, at subclass 2, the
forming regions of the broad, permitted lines are possibly obscured by dust clouds and/or
an unfavourable viewing angle.
M77 shows one of the larger red shifts of the Messier
galaxies. In the region of the Hα line it is about 24 Å
(ߣ ∙ ݖ). In the graph on the right, the scale and the blue
profile are based on the rest wavelength ߣ, calibrated
with known lines. The red profile is absolutely scaled
here with the calibration light; that on Table 73 refers
to the rest wavelength ߣ.
For obvious reasons the spectrum shows strong simi-
larity to profiles of galactic emission nebulae, which
are presented in sect. 28. Accordingly, the excitation
class ܧcan also be determined here. Already the He II
line at λ4686 shows, that ܧmust be > ܧ4. The crite-
rion log (ܫேଵାேଶ /ܫு ூூ (ସ଼)) results here in ~1.5, cor-
responding to an excitation class ܧ10 on the 12-step
Gurzadyan scale (sect. 28). This high excitation level is
also documented by the considerable intensity of the forbidden [O III] and [N II] lines, com-
pared to the rather weak H-Balmer series. So here even [N II] (6583) surpasses clearly the
Hα emission. Comparable to the supernova remnant M1 (sect. 28), the sulphur [S II] dou-
blet (λλ 6718/6733) is here also strikingly intense – probably as well due to shock wave-
induced excitation.
Veilleux and Osterbrock use in their classification scheme [307] inter alia the decadic loga-
rithms of the intensity ratios [ܱ (]ܫܫܫହ)/ ߚܪand [ܰ (]ܫܫହ଼ଷ)/ߙܪ. According to Shuder
and Osterbrock (1981), the intensity ratio
[ܱ (]ܫܫܫହ)/ > ߚܪ3
applies as a rough criterion to distinguish Seyfert- from H II- or Starburst galaxies (compare
with M82, Table 72!).
The direct vicinity of the supermassive black hole appears in the center of Seyfert galaxies
almost star like and very bright, while the rest of the galaxy remains relatively dark. This
considerably simplifies the recording of the spectrum of M77 with an apparent magnitude
of V 8.9m. The spectrum from the core was recorded with the 25μm slit and the 200L grat-
ing. Exposure time: C8/DADOS/Atik 314L+: 2x1200 sec, 2x2 Binning Mode, –20°C.
Spectroscopic Atlas for Amateur Astronomers 126
26.10 Quasars
The Quasar Phenomenon
The term "Quasar" is derived from Quasi-stellar Object (QSO),
because these objects appear as point shaped light sources.
Maarten Schmidt discovered the first in 1963 at the coordi-
nates of a corresponding entry in the mentioned radio source
catalogue (HST image: 3C273). It quickly became clear that
this object showed the largest Redshift, known at that time,
and therefore could impossibly be a star. In addition, the ob-
tained spectra differed dramatically from stellar profiles and
appeared more like those of Wolf Rayet stars, Nova outbursts,
or even Supernova explosions. According to today's knowledge, Quasars are considered as
the most energetic and luminous version of galaxies with active nuclei (AGN). The center of
such an object always hosts a supermassive Black Hole which accumulates matter from the
surrounding galaxy by an accretion disk and ejects a jet with relativistic speed. Therefore,
Quasars are also strong sources of X-ray and some of them also of radio emission ("radio
loud"). The mass estimation of the Black Hole is still difficult and uncertain. The literature
shows strongly scattering values for example [311], proposing a mass of some 1 bn. M .
Their point shaped appearance can be explained by the enormous brightness of the nuclei,
which in most cases totally outshine the rest of their host galaxies. Apart from the episodi-
cally occurring Supernova explosions they are by far the most luminous objects in the uni-
verse.
Table 75: Quasar 3C273, details of the spectral profile
Preliminary remarks: Here follows a summary of my publication “Quasar 3C273, Optical
Spectrum and Determination of the Redshift [36]. There, further information, as well as a
finder chart and details, about the recording of this object, can be found. The apparently
brightest Quasar 3C273 (mV ≈12.85m) in Virgo is often called the most distant object which
can still be seen with average amateur means, purely visually and without the use of as-
tronomy cameras. The designation means the object number 273 in Ryle's third Cambridge
catalogue of radio sources from 1959.
Heliocentric parameters according to NED [501] and other sources [1].
Redshift [501] Doppler Radial Relativistic Radial Distance
ߣ∆ = ݖ/ߣ velocity [501] velocity [500]* [bn. ly] [1]
ݒ = ܿ ∙ ݖ ( ݖ+ 1)ଶ − 1
ݒ = ܿ ∙ ( ݖ+ 1)ଶ + 1
+0.1583 +47’469 km s-1 +43‘751 km s-1 2.4
*Remark: The value for the relativistic radial velocity is currently no longer included in [500]
Spectral features of 3C273
The spectrum is dominated by extremely broadened emissions of the H-Balmer series and
by forbidden [OIII] lines at λλ 5007 and 4959, fusing here to a blend. Due to quantum me-
chanical reasons, the forbidden O III lines can't appear to be significantly broadened. It is
therefore discussed, whether the λ 5018 emission of the Fe II (42) multiplet (λλλ 4923,
5018 and 5169), supplies the major contribution to the intensity of this emission [309],
[312]. This multiplet frequently appears in the spectra of Active Galactic Nuclei (AGN), as
well as in the profiles of Protostars (sees PMS sect. 14). Striking is the much lower inten-
sity of the O III emission, relative to Hβ. This observation is in contrast to the spectral pro-
Spectroscopic Atlas for Amateur Astronomers 127
files of active Seyfert-type Galaxies (Table 73), planetary nebulae and H II regions (sect.
28). This phenomenon has already been noted by the discoverers of Quasars in the
1960ies. Even today just hypotheses exist about this issue.
Undisputed is the Ne III emission at λ 3869. The other features are mostly broad-band
blends of different emissions, generated by various ions. This significantly complicates the
line identification [312]. Consequently, their exact composition is still unclear. Striking is a
broad emission between λλ 4500 – 4700. J. B. Oke [313] suggested as the cause the He II
line at λ 4686 and numerous emissions of C III and N III - this in analogy with similar spec-
tra of Supernovae and Wolf Rayet stars. In [312] it is assumed that the striking emission in
the range of λ 5870 is caused by He I at λ 5876. Under discussion is also the Na I doublet,
which in certain phases can be observed during Nova eruptions. Due to the extreme shock
sensitivity of the metastable initial states of the forbidden [O III] lines, and the very low
ionisation energy of Na I, these emissions must necessarily be generated at a considerable
distance to the supermassive Black Hole. An indication for the contraction process within
the accretion disk are the inverse P Cygni profiles in the area around λλ 6100 - 6400, also
observable in the spectra of the much smaller disks around the T Tauri and Ae/Be- Proto-
stars (see sect. 14).
The Hα emission is redshifted here so far (1017 Å) that it coincides with the intense, tellu-
ric Fraunhofer A line. This is the cause why Hα appears split here [313]. This circumstance
complicates the determination of the Redshift, using this line, and seems at least to be one
of the reasons for the by far too low Balmer Decrement (ܫுఈ/ܫுఉ < 2.85) [313].
Because Hα appears split, the radial velocity in the vicinity of the Black Hole is estimated
here, using the FWHM of the Doppler-broadened and Gauss fitted Hβ line. Info to the for-
mula see [30], sect. 17. At such extreme line widths the correction of the instrumental
broadening can be neglected. It results in ܯܪܹܨா௦௦ ுఉ ≈ 88Å. The radial velocity of the
matter, calculated with the Doppler principle results in:
ݒ ≈ ܯܪܹܨா௦௦ ுఉ ∙ ܿ ≈ 5400 ݇݉/ݏ
ߣ ுఉ
This is roughly within the strongly scattering literature values. For the jet however, based
on X-ray analyses, up to 70% of the light speed are postulated [308]. Anyway this feature is
for amateurs, in the accessible optical spectral range, neither detectable nor measurable.
Not only the brightness of the object can vary considerably (see AAVSO), also the spectral
characteristics can change dramatically within a short time, such as the half-width ܹܧof
the Hβ emission (for details see [36]). It also shows that this central region can’t be indefi-
nitely large, and according to www.hubblesite.org can therefore hardly exceed the diameter
of our solar system. 3C273 would certainly be a highly interesting candidate for a monitor-
ing project. In addition to such considerations we should always be aware that these
changes, observed within a very short time, took place about 2.4 billion years ago, when
our earth was still in the Precambrian geological age!
Table 76: Quasar 3C273, Redshift
On this table, redshifted wavelengths are indicated, measured at Gaussian fitted lines and
related to the original wavelength scale, calibrated with a calibration light source [35]. The
amounts of redshift are not constant but, for a given radial velocity ݒ, proportional to the
rest wavelength ߣ of the corresponding spectral line. Therefore, the calculation of the ݖ-
values for all evaluated lines must yield almost the same amounts. This is also a good final
test for the quality of the wavelength calibration.
In the field of astrophysics for such highly red shifted objects, the distance is usually di-
rectly expressed as ݖ-value. It can easily be determined by the measured Redshift and re-
Spectroscopic Atlas for Amateur Astronomers 128
mains fully independent of assumptions for cosmological model parameters (eg Ω). Due to
the constant speed of light ܿ = ܿݐݏ݊, ݖis also used as a measure of time for the past.
∆ߣ
ߣ = ݖ
The ݖ-value, measured and calculated from this profile, amounts to = ݖ0.158 and is consis-
tent up to three decimal places with the literature value of +0.1583; for details see [30]
and [36]! The ∆ߣ values allow the calculation of the radial velocity ݒ with the usual Doppler
formula. With such high redshifts >1000 km/s, however, the relativistic formula (SRT)
should be used for ݒ [16].
∆ߣ ( ݖ+ 1)ଶ − 1
݈ܽݑ݉ݎ݂ ݎ݈݁ܦ: ݒ = ߣ ∙ ܿ ܿ݅ݐݏ݅ݒ݅ݐ݈ܽ݁ݎ: ݒ = ܿ ∙ ( ݖ+ 1)ଶ + 1
At redshifts of > ݖ0.1 increasingly dominates the cosmological expansion of the so-called
space-time lattice and the kinematic peculiar motion of galaxies plays virtually no more role
[30]. Because, however, the validity of both, the classical and the relativistic Doppler for-
mula, is limited to kinematic processes, their advanced application on the cosmologic space
expansion is currently being rejected by most of the experts. 3C273, with = ݖ0.158, is al-
ready something beyond this limit. Therefore in this case, both, the expansion velocity of
the space as well as the distance should be calculated with accordingly parameterised
cosmological models, mostly based on the ART. Anyway, if these values are calculated nev-
ertheless "conventionally", applying = ݖ0.158, it results for the "radial velocity":
ݒ = 47‘490km/s ݒ = 43‘808km/s
With this "radial velocity" and applying the Hubble law, finally the distance of some 2 billion
ly can be estimated (accepted value some 2.4 billion light years). Further details see [30]
and [36].
Spectroscopic Atlas for Amateur Astronomers 129
26.11 Blazars and BL Lacertae Objects (BL LAC's)
The Blazar Phenomenon
The term "Blazar" is composed of the terms "BL Lacertae" and "Quasar", used also abbrevi-
ated as "BL Lac". Allegedly the term was coined in 1978 by Ed Spiegel at a congress on BL
Lacertae objects in Pittsburgh, on the occasion of a banquet speech. Basically, these ob-
jects differ from the quasars just by our perspective, ie here we look directly into the syn-
chrotron radiation of the jet matter, which is ejected from the central black hole, parallel to
the rotation axis and almost with the speed of light.
Therefore, we observe here a highly intense, aperiodically fluctuating radiation with a
strong polarisation across the whole electromagnetic spectrum. Similar to the quasars,
these objects were first interpreted as blue variable stars. Not until 1968 their true nature
was first discovered at BL Lacertae. The spectra of these bright jets of matter typically
show no spectral lines, neither in emission nor in absorption. Therefore the red shift may
here only be measured at the comparatively very faint and diffusely appearing galaxy,
which is feasible just with large telescopes.
Table 77: Blazar Makarian 421 (Mrk 421)
Heliocentric parameters according to NED [501]
Radial- Redshift ݖDistance Type
Velocity. ݒ
+9'000km/s + 0.030 400 Mio Blazar/BL LAC
ly
With 400 million ly distance Makarian 421 is clearly the clos-
est object of the category Blazar/Quasar. It is still orbited by a
small companion galaxy Mrk 421-5 (details see [314]). Atlas
Image: courtesy of 2MASS/UMass/IPAC Caltech/NASA/NSF.
The apparent brightness of the Blazar is indicated by CDS
[500] with mV ≈ 12.9m, almost the same value as for the qua-
sar 3C273 (sect. 26.10). Anyway in most cases, the bright-
ness is significantly weaker here, ie in the range between 13m
and 14m (see AAVSO). Thus for amateurs this object is much
more difficult to record, but otherwise very easy to find, be-
cause it is located in an eye-catching star pattern just nearby
to 51 UMa (mV = 6.0m). With amateur equipment from Mrk 421, only a jet spectrum, lacking
of any spectral lines, can be recorded, which of course does not allow any determination of
the redshift. Table 77 shows this profile, which just shows the telluric absorptions. Re-
cording info: C8/DADOS, Grating 200L/mm, Atik 314L+ 1x1800sec, 2x2 binning.
Spectroscopic Atlas for Amateur Astronomers 130
26.12 List of Quasars Brighter than Magnitude 15m (DVAA)
The following table is an excerpt of the DVAA-List, containing Quasars brighter than Magni-
tude 18m [316]. The covered distance range here is 122 – 3'507 Mpc, z: 0.03 – 1.34.
Mag. Object Const. RA DEC Type z Dist.
[Mpc]
12.8 1ES 1959+650 DRA 19 59 59.9 +65 08 55 BL 0.05 183
12.9 3C 273 VIR 12 29 06.7 +02 03 08 QSO 0.16 584
12.9 MK 421 UMA 11 04 27.2 +38 12 32 BL 0.03 122
13.0 IRAS 01072-0348 CET 01 09 45.1 -03 32 33 QSO 0.05 210
13.1 MK 509 AQR 20 44 09.7 -10 43 24 QSO 0.04 137
13.7 MK 501 HER 16 53 52.2 +39 45 36 BL 0.03 130
13.8 MK 926 AQR 23 04 43.5 -08 41 08 QSO 0.05 183
13.9 HE 1029-1401 HYA 10 31 54.4 -14 16 52 QSO 0.09 329
13.9 MCG 11-19-006 UMI 15 19 21.6 +65 34 40 QSO 0.04 172
14.0 1ZW 1 PSC 00 53 34.9 +12 41 36 QSO 0.06 236
14.2 PG 1211+143 COM 12 14 17.7 +14 03 13 QSO 0.09 325
14.2 3C 371 DRA 18 06 50.7 +69 49 28 BL 0.05 199
14.2 KUV 18217+6419 DRA 18 21 57.3 +64 20 36 QSO 0.3 1035
14.3 PG 1351+640 DRA 13 53 15.7 +63 45 46 QSO 0.09 336
14.4 1E 0754+395 LYN 07 58 00.1 +39 20 30 QSO 0.1 365
14.4 MR 2251-178 AQR 22 54 05.9 -17 34 55 QSO 0.07 262
14.4 TON 599 UMA 11 59 31.9 +29 14 45 QSO 0.73 2203
14.4 HS 0624+6907 CAM 06 30 02.4 +69 05 04 QSO 0.37 1254
14.5 MK 180 DRA 11 36 26.5 +70 09 28 BL 0.05 180
14.5 TON 951 LYN 08 47 42.5 +34 45 05 QSO 0.06 247
14.5 IRAS 17596+4221 HER 18 01 09.1 +42 21 44 QSO 0.05 206
14.5 IRAS 21219-1757 CAP 21 24 41.7 -17 44 46 QSO 0.11 427
14.6 MK 478 BOO 14 42 07.5 +35 26 23 QSO 0.08 296
14.6 PG 1718+481 HER 17 19 38.4 +48 04 13 QSO 1.08 2996
14.6 7ZW 118 CAM 07 07 13.2 +64 35 59 QSO 0.08 303
14.6 MK 1298 LEO 11 29 16.7 -04 24 08 AGN 0.06 232
14.6 2ZW 136 PEG 21 32 27.8 +10 08 19 QSO 0.06 244
14.7 PG 1634+706 DRA 16 34 29.0 +70 31 33 QSO 1.34 3507
14.7 MK 304 PEG 22 17 12.2 +14 14 21 QSO 0.07 259
14.7 1E 2124-149 CAP 21 27 32.4 -14 46 48 QSO 0.06 221
14.7 1ES 0229+200 ARI 02 32 48.6 +20 17 17 BL 0.14 518
14.7 PG 0804+762 CAM 08 10 58.5 +76 02 43 QSO 0.1 380
14.7 TON 1388 LEO 11 19 08.8 +21 19 18 QSO 0.18 649
14.9 MS 03180-1937 ERI 03 20 21.2 -19 26 31 QSO 0.1 394
14.9 OF-109 ERI 04 07 48.5 -12 11 36 QSO 0.57 1815
14.9 MK 1383 VIR 14 29 06.6 +01 17 06 QSO 0.09 329
14.9 WAS 26 LEO 11 41 16.1 +21 56 22 QSO 0.06 244
14.9 MS 15198-0633 LIB 15 22 28.8 -06 44 41 QSO 0.08 322
14.9 MK 734 LEO 11 21 47.1 +11 44 19 AGN 0.05 191
15.0 PKS 0219-164 CET 02 22 00.8 -16 15 17 QSO 0.7 2128
15.2 HE 1211-1322 CRV 12 13 46.3 -13 38 52 QSO 1.13 3084
15.2 OP+313 CVN 13 10 28.7 +32 20 44 QSO 1 2813
Spectroscopic Atlas for Amateur Astronomers 131
TABLE 70 Richard Walker 2010/09©
Hα 6562.82
Comparison of Spectra: Galaxis M31 vs. Vindemiatrix G8 lllab Vindemiatrix
G8 lllab
Na l 5890/96
M31
Magnesium Triplet
5167 - 84
Hβ 4861.33
G-Band CH 4299 – 4313
Fraunhofer H + K
Spectroscopic Atlas for Amateur Astronomers 132
TABLE 71 LINER Galaxy M 94 NGC 4736
Telluric O2 6870 Scale: Rest-wavelength λ0 Redshift z = +0.0010
Two Profiles, recorded at different dates and with different slit positions
Hα 6562.82
Telluric O2 ©Richard Walker 2014/04
Na I 5889/95
Fe l/Ca l 5270
Mg l-Triplet
5167. - 5183
Hβ 4861.33
Hγ 4340.5
CH 4299 - 4313
Ca II 3968.47
Ca II 3933.66
Spectroscopic Atlas for Amateur Astronomers 133
Starburst Galaxy M 82 NGC 3034TABLE 72 [S II] 6717/31
Scale: Rest-Wavelength λ0[N II] 6583.6
Richard Walker 2013/04©Hα 6562.82
Hα 6562.82
Na I 5889/95
Hβ 4861.33
[O III] 5006.84
Hβ 4861.33
Hγ 4340.5
Spectroscopic Atlas for Amateur Astronomers 134
TABLE 73 [S II] 6717/31 Richard Walker 2013/04©
[O I] 6300.2
[N II] 6583.6
Hα 6562.82
[N II] 6548.1
Scale: rest wavelength λ0 Criterion log (IN1+N2 /IHe II (4686))≈ 1.5
Excitation class E10
Seyfert Galaxy M 77 NGC 1068
Classification lines[O III] 5006.84
[O III] 4958.91
Hβ 4861.33
He II 4685.7
Hγ 4340.5
[Ne III] 3967.47
[Ne III] 3868.76
Quasar 3C273 Line Identification Spectroscopic Atlas for Amateur Astronomers
Telluric Fraunhofer A
DADOS: Grating 200L mm-1, 50μm slit, recorded may 26, 2012 with Atik 314L+ -10°C, 5x1200s
The indication of the wavelength is provided in rest wavelength λ0 TABLE 75 Hα 6562.82
The profile is normalised to the continuum Ic = 1, the intensity on the level of the wavelength axis is Ic = 0.6
Na I 5889/95 [7] ?
I [O lll] 5006.84 / 4958.91 He I 5876 [7] ?
Fe II (42) 4923/5018/5169 Inverse P Cygni
1.0 Profiles
0.6 Hβ 4861.33
tellur. tellur.
He II 4686 [8] ?
C III / NIII [8] ?
Hγ 4340.47
Hδ 4101.74
Hε 3970.07
Ne III 3868.74
Red shifted, calibrated original scale 135
Rest wavelenght λ0
©Richard Walker 2013/04
Quasar 3C273, Redshift of the Hydrogen Balmer Lines Spectroscopic Atlas for Amateur Astronomers
DADOS: Grating 200L mm-1, 50μm slit, recorded may 26, 2012 with Atik 314L+ -10°C, 5x1200s TABLE 76
The indication of the wavelength, determined with Vspec at Gaussfits, is provided red shifted on the
original scale. The profile is normalised to the continuum Ic = 1, the intensity on the level of the
wavelength axis is Ic = 0.6.
Hα 7580
Δλ≈1017Å
I Hβ 5632
Δλ≈771Å
Hγ 5023
Δλ≈683 Å
Hδ 4748
Δλ≈646 Å
1.0
0.6 136
Red shifted, calibrated original scale
©Richard Walker 2013/04
Blazar Makarian 421 Spectroscopic Atlas for Amateur Astronomers
Scale: Rest-wavelength λ0 Redshift z = + 0.0300 Radial Velocity 9‘000km/s TABLE 77
©Richard Walker 2014/03Telluric O2 6870
137
Spectroscopic Atlas for Amateur Astronomers 138
27 Star Clusters
27.1 Short Introduction and Overview
Star clusters are formed by compacted stellar accumulations within or in the vicinity of gal-
axies. Most of the members of such clusters have evolved from a common gas cloud, and
have therefore about the same age. The individual clusters show a wide dispersion with re-
spect to density, age, and number of stars. Basically, star clusters can be divided in to the
following two main categories with completely different properties.
27.2 Open Clusters
Such clusters show an irregular shape and usually contain a few hundred (eg Pleiades) up
to at most a few thousand stars (eg, H and Χ Persei). The internal gravitational forces of
such clusters are usually too weak to hold together their stars for more than at most a few
hundred million years. One of the few exceptions is the very dense M67, whose age is es-
timated to about 4 billion years. The typical diameters of open star clusters are relatively
similar: eg Hyades: 15 ly, M67: 26 ly and Pleiades: 14 ly. Very small accumulations, just
consisting of a few stars and showing a common direction of movement, are referred as
"associations". Today, more than 1000 Open clusters are registered, which are distributed
over the entire visible region of the Milky Way. The observation of extragalactic open clus-
ters, eg in M31, remains reserved to large professional telescopes.
27.3 Globular Clusters
They differ from the open clusters in almost all aspects:
– they are strikingly spherical
– they are "packed" much denser
– they are all significantly older
– they contain up to several hundred thousand stars
– their typical diameter is more than 10 times larger
The typical distance between the stars in the outskirts of globular clusters is about 1 ly. In
the core area these interspaces can even shrink roughly to the diameter of the solar sys-
tem! Typical diameters of the comparatively much larger globular clusters are at
M2: 182 ly, M3: 180 ly, M5: 178 ly, M13: 145 ly.
In the center of individual globular clusters (eg, M15) intermediate-mass black holes of a
few thousand solar masses have been detected. This shows a certain similarity with dwarf
galaxies. Currently in the Milky Way about 150 globular clusters are known. They all orbit
the galactic center at a distance of about 130,000 light years, forming a halo this way.
Thus, these objects could even be referred as "extra-galactic" if they would not be tightly,
gravitationally bound to the Milky Way. However, it seems certain that their enormous high
age of about 12 billion years, is about the same as that of the entire Milky Way, ie slightly
younger than the universe with 13.7 billion. In contrast to the open clusters the high den-
sity and internal gravitational forces, allow the structure of globular clusters virtually an ar-
bitrarily long life. The observation of extragalactic globular clusters, eg in M31, remains re-
served to large professional telescopes.
27.4 Spectroscopic Analysis of Star Clusters
In contrast to the galaxies, in which amateurs are not able to record individual stars but
only composite spectra in the integrated light, here both, the open- and, with some restric-
Spectroscopic Atlas for Amateur Astronomers 139
tions, also the globular clusters, allow the analysis of the brighter individual stars. In the
professional field, multichannel spectrographs allow the simultaneous recording of up to
several 100 profiles. Main objective here is mostly the determination of the metal abun-
dance ܼ "[݁ܨ/( ]ܪsect. 4.7). This value allows direct conclusions about the age of a star
cluster. The stars of the first generation, somewhat oddly called Population II, were created
with the birth of the Milky Way 12 billion years ago, where the interstellar matter was still
dominated by hydrogen, helium and lithium. The enrichment with heavier elements took
place only later by matter from SN explosions or repelled planetary nebulae. This enriched
material generated later on the metal-rich second star generation, similarly confusing
called Population I, to which belongs also our Sun. The most efficient way for the determi-
nation of the ܼ- value, has proven the analysis of the Ca II calcium triplet in the near infra-
red at λλλ 8542, 8498 and 8662. This process is abbreviated called as CaT. The empirical
relationship between the metallicity ܼ and the summed EW- values of the mentioned Ca II
absorptions has been refined over the last 25 years and outlined in numerous publications,
such as [325 – 327].
27.5 Spectroscopic Age-Estimation of Star Clusters by Amateurs
The described CaT-method is too demanding for most amateurs. So, for example, the Ca II
calcium triplet is located in the infrared range, somewhat outside the reach of most ama-
teur equipment and frequently contaminated by blends with other absorptions. In [30] sect.
14.6, a simple method is presented, based on the correlation between the duration of stay
on the main sequence MS and the spectral class of a star. It is based on:
– The assumption, that the individual stars of such clusters have been formed at about the
same time from a common gas- and dust cloud
– The known relation: The earlier classified (or more massive) the star, the shorter the
lifetime
– The quintessence: Within a stellar sample, the duration of the stay on MS for the earli-
est spectral class determines very roughly the age of the cluster.
At the so called Turn off Point the star reaches on the MS its ear- Spectral- Stay on MS
liest possible spectral class. Then it moves within the HRD to the class [Years]
top right of the giant branch, showing thereby significantly later O7 6 Mio.
classifications. The table to the right shows, how the stay of later O9 8 Mio.
classified stars on the MS increases dramatically. B0 12 Mio.
B1 16 Mio.
The table to the right shows that for later classified stars the stay B2 26 Mio.
on the MS increases dramatically [405]. Starting from the late G- B4 43 Mio.
class it even exceeds significantly the present age of the Universe B6 95 Mio.
of 13.7 billion years. Such data are a rough guide only. Based on A0 350 Mio.
model calculations they often show (source-dependent) a consid- A5 1.1 bn.
erable spread. F2 2.7 bn.
G2 9.4 bn.
Example: If within a cluster a larger spectral sample yields an K0 23 bn.
early A-type as the earliest classification, its age can roughly be
estimated to 350 – 500 million years. Any earlier classified MS
stars of types B and O, are much more short-lived and thus either
already exploded in a SN or migrated in the HRD to the top right
of the giant branch (see [30], sect. 14.6).
In professional fields, such studies are often carried out photometrically in the ܸ –ܤsystem.
The major disadvantage is that these measured magnitude values always appear reddened
by interstellar matter and need first of all to be adjusted with appropriate models. In the ta-
ble of sect. 38, the assignment of "de-reddened" ܸ –ܤmagnitudes to the corresponding
Spectroscopic Atlas for Amateur Astronomers 140
spectral class can be found. Anyway as the first option for amateurs remains clearly the
spectroscopic age-estimation based on the spectral class.
27.6 The Pleiades - Analysis by Individual Spectra
The approximately 390 light years distant Pleiades (M45) are so nearby, that this object
appears more as a small constellation, rather than an open cluster. With the unaided eye,
we see here just the brightest 6 up to max. 10 stars which astonishingly belong all to the
middle to late B-Class. Their visibility depends mainly on the seeing and the current bright-
ness of the Be-star Pleione.
Star Flamsteed HD No. Apparent Abs. [*] Spectr. Remarks (CDS)
name No. brightn. brightn. Class
Alcyone 25 Tau (η) 23630 ~2.87m –2.74 B7 IIIe Be- and Multiple star
–2.00 B8 III Binary star system
Atlas 27 Tau 23850 ~3.63m –1.84 B6 IIIe Be Star
–1.83 B8 III Binary star system
Electra 17 Tau 23302 ~3.7m –1.58 B6 IVe Variable β Cep Type
–1.31 B6 IV Binary star system
Maia 20 Tau 23408 ~3.87m –0.58 B8 V ne Be Star
–0.28 B7 IV Variable
Merope 23 Tau 23480 ~4.18m + 0.01 B8 V
+ 0.05 B8 V Variable
Taygeta 19 Tau 23338 ~4.3m
Pleione 28 Tau 23862 ~5.09m
Celaeno 16 Tau 23288 ~5.46m
18 Tau 23324 ~5.64m
Asterope 21 Tau 23432 ~5.76m
[*] Abt/Levato [328], other data CDS [500]
Image: Pleiades labelled with Flamsteed numbers (M. Huwiler)
28
27 18
25 21
20
19
23 16
17
Spectroscopic Atlas for Amateur Astronomers 141
Overall, the cluster contains about 500 stars. According to a study by Abt and Levato [328],
the brightest 50 stars of M45 spread over the various spectral classes as follows:
B: 17, A: 24 and F: 9. The remaining stars are classified as type F or later. Anyway it is still
an enigma, why in this rather small cluster the number of massive stars, classified in the
range of the middle to late B-Class, exceeds by orders of magnitude their expected average
incidence of just some 0.125% (see sect. 6.4)! The famous reflection nebulae are not, as
originally believed, the remaining "birth shells" of the brighter stars but rather a galactic gas
cloud, which is randomly located in the trajectory of M45.
Tables 78A and B:
Both tables impressively show the very similar spectra of the 10 brightest stars within
M45. The middle to late B class can be recognised here mainly due to the relatively intense
Balmer lines in combination with the very weak Ca II absorptions at λ 3933.66 (sect. 10).
Striking features of all profiles are also the small emission humps in the continuum at λλ ~
5840 and ~ 5350. Further visible here is the increase in intensity of the H-Balmer lines be-
tween the classes B6 to B8.
Table 78A shows a montage of those four spectra, which are classified to show emission
lines (index e). Three of them are even classified as Be-stars (CDS). However in March
2014, just two of them showed intensive emission lines and Merope, 23 Tau, a very weak
filling-in in the core of its Hα absorption. As the only one the Be- star Pleione shows in its
profile coarser absorptions between the H-Balmer lines. Recording info: C8/DADOS, Grat-
ing 200L/mm, Atik 314L+, average 1x60sec, 2x2 binning.
27.7 Age estimation of M45
The earliest main sequence (MS) classification is here B8. All earlier B types have already
migrated to the giant branch. According to the table in sect. 27.5 this yields an age of some
more than 100 million years. The accepted value is in the range of 100 – 130 million years.
Cecilia Payne Gaposchkin (see sect. 7) discovered in M45 several White Dwarfs. This
would mean that some stars of the middle to late B-type with < 8M , have already repelled
Planetary nebulae. Under debate is the hypothesis, if in exceptional cases, also early B stars
with > 8M , could end up as White Dwarfs. If such stars would lose sufficient matter, dur-
ing the Giant- or Post AGB Phase, their mass could sink this way below the Supernova limit.
27.8 Globular Clusters – Analysis by Integrated Spectra
Due to their enormous distances of about 25,000 – 40,000 ly the brightest individual stars
of globular clusters reach at most an apparent magnitude of mv ≈11m (eg. at M5 and M13).
Thus for an individual star analysis with an 8-inch telescope only the few brightest speci-
mens could be recorded with a slit spectrograph. Therefore recommendable for amateurs is
the recording of the composite spectrum in the integrated light, as it is applied also for the
galaxies in sect 26. Due to the very old stellar inventory almost exclusively absorption lines
can be seen here. With such spectra a so-called "Integrated spectral class" [330] can be
determined, and this way, by analogy with the open clusters, the age of the cluster can be
estimated.
In the professional field [331], these profiles are also compared with synthetic model spec-
tra and the Turn Off Point is determined, inter alia, by means of the Balmer lines. According
to [331] the integrated light-spectra are obtained here by sweeping over the cluster to pre-
vent, that the profile is determined by the stars of just one particular cluster zone. Due to
the relatively short focal lengths and the relatively inaccurate autoguiding, such an action is
not necessary at the amateur level. Anyway when recording the spectra in Table 79, the
image was slightly defocused in order to avoid any disproportionate influence of individual
stars.
Spectroscopic Atlas for Amateur Astronomers 142
A still unsolved enigma is posed by single blue stars of earlier spectral type, the so-called
"Blue Stragglers" (BSS). Such short-lived objects can be detected mainly in the central re-
gions of all known globular clusters and don't fit at all into the picture of an extremely old
cluster. There are discussed several hypotheses – one of which is that due to the high den-
sity of the central cluster zones, such "blue stragglers" could be generated by the fusion of
two or more red giants.
Table 79:
This table impressively shows the very similar integrated spectra of M3, M5 and M13. Not
identifiable remain here, in the profile of M13, just the intensive absorptions at ~λ 6600
(artefacts?). Well defined and recognizable are here the Balmer lines, the two Fraunhofer
Ca II absorptions, as well as, surprisingly intense and well defined, the molecular CH band.
Very easy is here the rough determination of the integrated spectral type with about F6 -
F7. The striking "brand" of the F-class is the combined appearance of the CH absorption at λ
4299 – 4313 and the directly adjacent Hγ line (sect. 16). The decimal subclass is derived
here from the nearly equal intensity of these two absorptions. The applied exposure time
with the C8 was here 1200s in the 2x2 binning mode.
27.9 Age Estimation of M3, M5 and M13
Applying this rather simple method with the middle to late F-class, it results here a some-
what too young age for the clusters of just about 6-7 bn. years (sect. 27.5). With 12 bn.
years the accepted value is clearly higher.
M3: Image HST Central part of M5: Image HST
Central part of M13: Image HST
Open Cluster M45 Pleiades Bright Stars with emission line classification Spectroscopic Atlas for Amateur Astronomers
Spectrograph: DADOS, Grating 200L/mm, Slit: 50µm TABLE 78A
Hβ 4861.33 Pleione 28 Tau B8 Vne Hα 6562.82
Alcyone 25 Tau B7 IIIe Telluric O2
Hγ 4340.47 Electra 17 Tau B6 IIIe
Merope 23 Tau B6 IVe Na I 5890/96
Hδ 4101.74 Interstellar
He I 4025.5
Hε 3970.07
Ca II 3933.66
H8 3889.05
H9 3835.38
143
©Richard Walker 2014/04
Open Cluster M45 Pleiades Bright Stars with absorption-lines Spectroscopic Atlas for Amateur Astronomers
Spectrograph: DADOS, Grating 200L/mm, Slit: 50µm TABLE 78B
Hβ 4861.33 Taygeta 19 Tau B6 IV Hα 6562.82
Celaeno 16 Tau B7 IV
Hγ 4340.47 Atlas 27 Tau B8 III Telluric O2
Maia 20 Tau B8 III
Hδ 4101.74 Asterope 21 Tau B8 V
He I 4025.5 18 Tau B8 V
Hε 3970.07
Ca II 3933.66
H8 3889.05
H9 3835.38
144
©Richard Walker 2014/04
Spectroscopic Atlas for Amateur Astronomers 145
TAFEL 79 ??
Hα 6562.82 ©Richard Walker 2014/04
Na I 5890/96
Hβ 4861.33Globular Clusters M3, M5, M13
Hγ 4340.47Spectrograph: DADOS, Grating 200L/mm, Slit: 50µm
CH 4299 - 4313
M5 NGC 5904
Hδ 4101.74 M3 NGC 5272
Ca II 3968.47 M13 NGC 6205
Ca II 3933.66
Spectroscopic Atlas for Amateur Astronomers 146
28 Emission Nebulae
28.1 Short Introduction and Overview
Reflection nebulae are interstellar gas and dust clouds which passively reflect the light of
the embedded stars. Emission nebulae however are shining actively. This process requires
that the atoms are first ionised by hot radiation sources with at least 25,000K. The density
of the nebulae is so extremely small that on earth it can be generated only as the best ultra-
high vacuum. In [30], sect. 22, these processes are explained more in detail. The require-
ments for the development of emission nebulae are mainly met by the following astronomi-
cal “object classes”.
28.2 H ll Regions
Textbook example is the Orion Nebula M42 (Photo: NASA).
Here extremely hot stars of the O- and early B class ionise –
in addition to helium, oxygen and nitrogen – primarily hydro-
gen atoms of the surrounding nebula. This requires UV pho-
tons, above the so-called Lyman limit of 912 Å and corre-
sponding to an ionisation energy of >13.6 eV. This level is
only achievable by very hot stars of the O- and early B-Class.
Such H II regions tend to have a clumpy and chaotic structure
and may extend over dozens of light years. They show a high
star formation rate and can still be detected even in distant
galaxies. The reddish hue is caused by the dominant Hα emission.
28.3 Planetary Nebulae PN – The Most Significant Subgroup of Emission Nebulae
In the central part of these much higher energetic objects are
mostly extremely hot white dwarfs with up to > 200,000K.
This is the final stage of stars at the end of the AGB (sect. 20
– 23) with originally <8 solar masses. They ionise the atoms
of their rather slowly expanding former stellar envelopes
(some 20-40 km/s). Photo: NGC 6543 Cat’s eye nebula
(NASA). About 10% of central stars show similar spectra in
the final stage like Wolf Rayet stars, and thus have WR-
classifications (WRPN).Their absolute magnitude however is
considerably lower. PN often show an ellipsoidal shape, in
some cases with a regular fine structure. The reasons for the
numerous other existing forms are only partially understood.
28.4 Protoplanetary Nebulae
They form the precursors of planetary nebulae. They are
excited by the so-called post AGB stars – former carbon
stars, ie Mira variables at the upper end of the asymptotic
giant branch AGB, just beginning to repel their envelops
(sect. 23). They are still not hot enough to excite higher ion-
ised emission lines such as [O III]. Textbook example is the
Red Rectangle Nebula in the constellation Hare (Table 85).
Spectroscopic Atlas for Amateur Astronomers 147
28.5 Supernova Remnants SNR
SNR show a striking filamentary structure. The main part of
the ionisation energy is provided here by the collision of
the rapidly expanding stellar envelope (a few 1000 km/s)
with the interstellar matter. Photo: M1 Crab Nebula
(NASA). In the center of SNR the remaining Neutron Star or
Pulsar emits a wind of relativistic electrons with nearly
light speed. It is deflected or slowed by magnetic fields
within the plasma or electric fields around the ions. Such
energy transformations are compensated by emitted pho-
tons, causing broadband Synchrotron- or Bremsstrahlung,
predominantly in the X-ray domain.
28.6 Wolf Rayet Nebulae WR
The shells around the high energy Wolf-Rayet Stars are ex-
cited in a similar way, but anyway much less intensive,
such as the Crescent Nebula NGC 6888 (Table 87) or
Thors Helmet NGC 2359 (Table 88). The picture on the
right, by ESO, shows the Carina Nebula NGC 3372 with the
ionising source WR 22.
28.7 Common Spectral Characteristics of Emission Nebulae
Besides of the chemical composition the local state of the plasma is determined by the
power of the UV radiation as well as the temperature ܶ and density ܰ of the free elec-
trons. By recombination the ions recapture free electrons which subsequently cascade
down to lower levels, emitting photons of well defined discrete frequencies (fluorescence
effect). According to the Plank formula: ∆ = ܧℎ ∙ ߥ, the frequency of the emitted photon ߥ
[nu:] is exactly proportional to the energy difference ∆ ܧbetween the levels, of the down-
ward electron transition. For these reasons, emission nebulae generate, similar to a gas
discharge lamp, predominantly “quasi monochromatic” light, not as a continuum, but rather
as several discrete emission lines. Accordingly effective are therefore specifically designed,
narrow band nebula filters. With the exception of SNR, emission nebulae show only a very
faint continuum radiation.
Since the main part of the light is concentrated on a few, more or less intense emission
lines, these objects can be still detected even at extreme distances. The brightest [O III]
lines become photographically visible just after very short exposure times. In all types of
emission nebulae physically the same ionisation processes are responsible for the line for-
mation, albeit with very different excitation energies. This explains the very similar appear-
ance of the spectra. The graph below shows a section of the emission spectrum of M42
with two noticeable features:
Spectroscopic Atlas for Amateur Astronomers 148
1. The intensity ratio of the brightest [O III] lines is always: (ܫ5007)/(ܫ4959) ≈ 3.Olll 5006.84Hα 6562.82
2. The intensity ratio between the hydrogen lines, called Balmer Decrement D, representsOlll 4958.91
the quantum-mechanically induced intensity loss of these lines in the direction of decreas-Hβ 4861.33
ing wavelength. Detailed description see [30], sect. 20.
Important for astrophysics is the intensity
ratio )ߙܪ(ܫ = ܦ/)ߚܪ(ܫ. The theoretically cal-
culated value for thin gases is ்ܦ ≈ 2.85.
The steeper the curve, the greater is the se-
lective interstellar extinction (reddening) of
light by dust particles, what in Photometry is
called red colour excess. As a result, the
lines at shorter wavelengths are increasingly
shown too short. Most of the galactic PN,
reachable for amateurs, show values of
≈ ܦ3.0 − 3.3 [203]. Therefore, for a rough
determination of the excitation class this ef-
fect can be neglected, especially as the di-
agnostic lines are relatively close together (see below). However, there are stark outliers
like NGC 7027 with D ≈ 7.4 [14]. For extragalactic objects ܦbecomes > 4, which in any
case requires a correction of the intensities ("Dereddening") [204].
28.8 Emission Line Diagnostics and Excitation Classes ࡱ
Since the beginning of the 20th Century numerous methods have been proposed to deter-
mine the excitation classes of emission nebulae. The 12-level “revised” Gurzadyan system
[14], which has been developed also by, Aller, Webster, Acker and others [204, 205, 206]
is one of the currently best accepted and appropriate also for amateurs. It relies on the
simple principle that with increasing excitation class, the intensity of the forbidden [O III]
lines becomes stronger, compared with the H-Balmer series. Therefore as a classification
criterion the intensity sum of the two brightest [O III] lines, relative to the Hβ emission, is
used. Within the range of the low excitation classes E: 1–4, this value increases strikingly.
The [O III] lines at λλ4959 and 5007 are denoted in the formulas as ܰଵ and ܰଶ.
For low Excitation Classes E1 – E4: ܫேଵାேଶ /ܫுఉ
Within the transition class E4 the He II line at λ4686 appears for the first time. It requires
24.6 eV for the ionisation, corresponding to about 50,000K [202]. That's almost twice the
energy as needed for H II with 13.6 eV. From here on, the intensity of He II increases con-
tinuously and replaces the now stagnant Hβ emission as a comparison value in the formula.
The ratio is expressed here logarithmically (base 10) in order to limit the range of values for
the classification system:
For middle and high Excitation Classes E4 – E12: log(ܫேଵାேଶ /ܫு ூூ (ସ଼))
The 12 ܧ-Classes are subdivided in to the groups Low ( = ܧ1 − 4), Middle ( = ܧ4 − 8) and
High ( = ܧ8 − 12). In extreme cases 12+ is assigned.
Spectroscopic Atlas for Amateur Astronomers 149
Low Middle High
–ܧClass –ܧClass log(ܫேଵାேଶ /ܫସ଼)
ܫேଵାேଶ /ܫுఉ –ܧClass log(ܫேଵାேଶ /ܫସ଼)
E1 0–5 E4 2.6 E9 1.7
E2 5 – 10 E5 2.5
E3 E6 2.3 E10 1.5
E4 10 – 15 E7 2.1 E11 1.2
>15 E8 1.9
E12 0.9
E12+ 0.6
28.9 Remarks to the Determination of Excitation Classes and Recording of Spectra
The determination of the low E classes 1–4 is easy, since the Hβ line, compared to the
[O III] emission, is relatively intense. At level E4 the He II line (λ 4686) starts very weak, re-
quiring very low-noise spectra, and a strong zoom into the intensity axis.
Quite easy to record are spectra from the very small, disc-shaped and blue-greenish shining
PN. Thus they are very quickly to localize within a stellar group and the exposure time for
the bright representatives takes only a few minutes (200L reflection grating). The brightest
[O III] line appears often just after a few seconds on the screen (eg NGC 6210). The inten-
sity of the lines is here integrated along the very short, exposed part of the slit. But along
this very short appearing diameter the individual lines show considera-
bly different intensities. Furthermore, during long exposure times, small
changes in the slit position with respect to the nebula are to observe as
a result of inadequate seeing and/or guiding quality. Tests have shown
that several spectra of the same object may show significantly different
results including for the Balmer Decrement. Anyway the influence on
the excitation class was observed as quite low! The picture shows the
small sliver of the Spirograph Nebula (IC 418) on the 25 μm slit (PHD
Guiding). Between the green autoguiding cross and the slit the bright
central star is visible.
By contrast the large appearing nebulae as M27 and M57 require, with the C8 and the
Meade DSI III, at least 20–30 minutes of exposure time (without binning) and a totally
cloud- and haze-free section of the sky. But they allow a selective recording of spectra
within specific areas of the nebula, and to gain an intensity curve along the quite long ex-
posed part of the slit. In this case, tests with several spectra of the same object showed
consistent results. More detailed information on this topic, as well as remarks to the correc-
tion of the raw profiles, can be found in [30] sect. 20 and 21.
Spectroscopic Atlas for Amateur Astronomers 150
28.10 The Excitation Class as an Indicator for Plasma Diagnostics
Gurzadyan (among others) has shown that the excitation classes are more or less closely
linked to the evolution of the PN [14], [206]. The study with a sample of 142 PN showed
that the E-Class is a rough indicator for the following parameters; however in the reality the
values may scatter considerably [13].
1. The age of the PN
Typically PN start on the lowest E- level and subsequently step up the entire scale with
increasing age. The four lowest classes are usually passed very quickly. Later on this
pace decreases dramatically. The entire process takes finally about 10,000 to > 20,000
years, an extremely short period, compared with the total lifetime of a star!
2. The Temperature ܶ of the central star
The temperature of the central star also rises with the increasing E-Class. By repelling
the shell, increasingly deeper and thus hotter layers of the star become "exposed". At
about E7 in most cases an extremely hot White Dwarf remains, generating a WR-like
spectrum. This demonstrates impressively the table of the PN in sect. 35. Hence, for
ܶ [K] the following rough estimates can be derived:
E-Class E1-2 E3 E4 E5 E7 E8-12
ܶ [ ]ܭ35,000 50,000 70,000 80,000 90,000 100,000 – 200,000
3. The Expansion of the Nebula
The visibility limit of expanding PN lies at a maximum radius of about 1.6 ly (0.5parsec),
because from here on the dilution becomes too great [202]. With increasing E-class,
also the radius of the expanding nebula is growing. Gurzadyan [206] provides mean val-
ues for ܴ [ly] which however may scatter considerably for the individual nebulae.
E-Class E1 E3 E5 E7 E9 E11 E12+
ܴ [݈ ]ݕ0.5 0.65 0.72 1.0 1.2 1.4 1.6
28.11 Emission Lines identified in the Spectra of Nebulae
The appearance and intensity of emission lines in the spectra of the individual Nebulae are
different. Therefore here follows a compilation with identified emission lines from Plasma
Recombination Lasers in Stellar Atmospheres [200] and Frank Gieseking [202]. So-called
"Forbidden lines" are written within brackets [].
Ne III 3869 [Ne III] 3967.5 He I 4026.2 [S II] 4068.6 Hδ 4101.7 C II 4267.3 Hγ 4340.5
[O III] 4363.2 He I 4387.9 He I 4471.5 He II 4541.6 [Mg I] 4571.1 [N III] 4641] He II 4685.7
[Ar IV] 4740.3 Hβ 4861.3 He I 4921.9 [Olll] 4958.9 [Olll] 5006.8 N l 5198.5 He II 5411.5
[Cl lII] 5517.2 [Cl lII] 5537.7 [O I] 5577.4 [N II] 5754.8 He I 5875.6 [O I] 6300.2 [S III] 6310.2
[O I] 6363.9 [Ar V] 6434.9 [N II] 6548.1 Hα 6562.8 [N II] 6583.6 He I 6678.1 [S II] 6717.0
[S II] 6731.3 [He II] 6890.7 [Ar V] 7006.3 He I 7065.2 [Ar III] 7135.8 He II 7177.5 [Ar IV] 7236.0
[Ar IV] 7263.3 He l 7281.3
28.12 Commented Spectra
Because spectra of emission nebulae barely show a continuum, the profiles in the following
tables are slightly shifted upwards to improve the visibility of the scaled wavelength axis.
The presentation of the following objects is sorted according to ascending excitation
classes. For most of the PN, the problem of distance estimation is still not really solved. The
information may therefore vary, depending on the sources up to >100%! Correspondingly
inaccurate are therefore also the estimated diameters of the nebulae!
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