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Survey on Gaseous Nebula. Reporter: Jinjuan Liu 2004.12.29. Intense Diffuse Far-UV Emission from the Orion Nebula. Jayant Murthy David J. Sahnow and R. C. Henry. Observations and Data Analysis. FUSE was launched on June 24, 1999 into a low Earth orbit (LEO) by a Delta II rocket

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survey on gaseous nebula

Survey on Gaseous Nebula

Reporter: Jinjuan Liu


intense diffuse far uv emission from the orion nebula

Intense Diffuse Far-UV Emission from the Orion Nebula

Jayant Murthy David J. Sahnow and R. C. Henry

observations and data analysis
Observations and Data Analysis
  • FUSE was launched on June 24, 1999 into a low Earth orbit (LEO) by a Delta II rocket
  • Two of the observations in the S405/505 program, were of a region of empty space near the star HD 36981 (Fig. 1, Table 1).

Fig. 1.— An interstellar dust map (Schlegel et al. 1998), with the Orion Nebula

(M42) in the center of the image.

Overplotted are the hot stars in Orionthe,

the star HD36981 :red circle at the NE edge of M42

The FUSE spectrographs are the first instruments with sufficient sensitivity to perform absorption line spectroscopy of the diffuse ISM
  • The diffuse emission is due to scattered light from the nearby stars
  • and a comparison of the spectra can provide important clues about the source of the photons and the geometry of the dust
our first assumption was that we were observing scattering of the light of HD 36981
  • But a comparison of the spectra near Ly shows that HD 36981 cannot be the major contributor to the scattered radiation ---- the broad intrinsic Ly absorption in the star is not reflected in our observed spectrum(fig.3)
  • Also Ori C (HD 37022;O5V; V = 5.1) — the brightest of the Trapezium stars — may contribute as much or more energy as HD 36981
  • Figure 2:The shape matches the diffuse spectrum well until about 1150 A when interstellar Ly absorption eats into the stellar spectrum.

Fig. 3.— Plotted are (from top to bottom) the FUSE spectrum of HD 36981, the Copernicus spectrum of Ori C and the FUSE observation reported here.


Fig. 2.— The spectrum for our diffuse observation is shown. Superimposed on the

plot is a scaled spectrum of Ori (red line) taken from the Copernicus archive.

The spectra shape is similar to the diffuse spectrum until the onset of interstellar Ly

absorption near 1200 °A.

Fig 3:
  • The Ly line is much broader in both stellar spectra than in the diffuse observation;

--- in the case of HD 36981 due to self-absorption in the star and in the case of Ori C due to interstellar absorption in the several interstellar clouds in front of the Trapezium

  • The Ly emission line is seen in the middle of the interstellar Ly (1026 A) absorption line in our diffuse FUSE observation.
  • Note that, surprisingly, the diffuse spectrum is rich in absorption lines which are not seen in either of the stellar spectra

---This is due to molecular hydrogen in both absorption and emission

it seems likely that we are observing the scattered light from Ori C through a much lower column density line of sight than the direct line to the star.
  • We have begun the process of modeling the entire spectrum in order to determine the scattering properties of the dust grains near Orion.
The authors believe that such high resolution observations of the diffuse scattering provide an important new tool for the study of interstellar dust scattering

--- in understanding the spectral properties of the interstellar dust

--- also in providing clues to the origin of

the diffuse emission.

far ultraviolet spectroscopic explorer observations of nebular o vi emission from ngc 6543


Robert A. Gruendl and You-Hua ChuandMartı´n A. Guerrero

observed: NGC 6543

with FUSE(Far Ultraviolet Spectroscopic Explorer)

to search for Nebular O vi emission.

Three observations of NGC 6543 were obtained using the HIRS aperture
  • “N-Cav” samples the northern edge of the central cavity,
  • “N-Ext” the northern extension of the cavity
  • “OFF” :a region outside the nebula at 150” north of the central star
NCav and N-Ext positions show prominent narrow emission lines of C ii ll1036, 1037, broad emission lines of O vi ll1032,1038
The continuum emission in the nebular spectra is most likely stellar emission scattered by dust in the nebular gas.
  • Figure 3: the continuum contamination is effectively removed
  • The O vi 1038 line flux and shape should be regarded with some skepticism
  • because: the emission line is close to a strong C ii absorption line and is in a spectral region where the P Cygni profile from the central star may be oversubtracted

Fig.3 The bottom left panel shows the echellogram, and the bottom right shows the line profile extracted from the position marked by the horizontal dashed lines on the echellogram. The systemic velocity of NGC 6543 is marked by a dashed vertical line in each panel.

after convert from the observed to the intrinsic intensity of the O vi emission from NGC 6543,
  • The intrinsic O vi line intensity is therefore times the observed intensity.
  • In Figure 3, we further compare the velocity profile of the O vi lines to that of the [O iii] 5007 line from an east-west–oriented long-slit echelle spectroscopic observation taken at a position 4”.5 north of the central star.
  • The brightest emission component in the [O iii] line image has low velocities and corresponds to the dense envelope of the nebular shell
  • the faint expanding “blister” near the center of the line image corresponds to the northern extension of the central cavity.
  • The heliocentric velocity of the O vi emission is similar to that of the expanding blister in the [O iii] line, indicating that they are physically associated
  • Since the central star of NGC 6543 is too cool to photoionize O v, the O vi ions must be produced by thermal collisions in gas at the interface layer between the cool nebular shell and the hot gas in the PN interior
  • ----Are the FUSE observations of O vi from NGC 6543 consistent with that expected from thermal conduction in such an interface layer?
  • The first model:
  • Assume:
  • PN shell is a pressure-driven bubble and the radiative losses of the hot interior gas have a negligible effect on the dynamics
  • a constant AGB mass loss
  • Solve the dynamics of such a bubble analytically ----characterized by a pressure and age
by numerically solving the continuity and energy equations----- radiative losses and thermal conduction are considered----calculate the temperature and density structure at the interface layer
  • Use these observational parameter of the nebula----we find the expected O vi emission to be ~10 times higher than observed
  • ----from the incorrect model prediction of pressure within the central cavity and incorrect assumptions about the fast and AGB wind mass-loss rate
Further consideration:
  • To obtain an alternative estimate of the O vi emission, we adopt the observed pressure obtained for the hot gas in the continuity and energy equations
  • to determine the temperature and density structure of the interface layer.
  • in Figure 4 ----The resulting temperature and density profile applicable to the N-Cav observations are plotted as a function of distance from the inner wall of the nebular shell.

Fig. 4.—Results of the model calculation for the N-Cav observation of NGC6543 showing profiles for the temperature (top), electron density (middle), and O vi emissivity (bottom) as a function of distance from the inner wall of the nebular shell.

Based on the calculations, we would expect and ergs cm-2 s-1 at the N-Cav and N-Ext positions, respectively.
  • While the predicted O vi emission intensity is roughly 1.5 times higher than observed,
  • the two are consistent
a close binary nucleus in the most oxygen poor planetary nebula pn g135 9 55 9
  • G. H. Tovmassian, R. Napiwotzki, M. G. Richer, G. Stasin´ska, A. W. Fullerton

And T. Rauch

  • observation of PN G135.9+55.9 (ID C034) was conducted through the LWRS aperture of FUSE on 2002 January 30 for a total exposure time of 30,660 s.
  • FUSE: the Far Ultraviolet Spectroscopic Explorer
  • The observations at the Canada-France-Hawaii Telescope(CFHT) were obtained on 2003 May 1 UT with the MOS spectrograph
  • The data were reduced using the Image Reduction and Analysis Facility ( IRAF) software package
the main features of the observed spectra
  • 1. the night exposure (NE) spectrum
  • No emission lines could be seen except those that are definitely associated with terrestrial airglow.
  • Most of these lines disappear in the NE spectrum with only the strongest lines of H i and O i leaving residual emission peaks whose intensities are dramatically reduced compared to the all-photon spectrum.

Fig. 1.—FUSE composite spectrum of PN G135.9+55.9 for the night-only exposure. A few of the strongest H i and O i k1025.76 airglow lines remain in

emission. The other features are absorption lines, most of which are due to interstellar matter. Molecular hydrogen lines are marked along the bottom of the

spectrum: the tick marks for the Lyman series lines are above those for the Werner series. Some other prominent interstellar lines of other elements are marked above

the spectrum. The rest positions of Ly and Ly are marked above the corresponding boxes.

The best opportunity for parameter estimation is offered by the hydrogen Balmer lines in the optical spectrum
  • it was possible to fit five Balmer lines (from Hγ to H9)
  • The model atmosphere analysis was done with the program FITSB2, Fit results are stellar parameters and RVS(radial velocities)
  • Because of the nebular contamination of the available spectra, it was not possible to determine both the temperature and gravity from the available spectra.
we kept the temperature fixed during the fitting procedure and determined only the surface gravity
  • The results are listed in Table 2.
  • combined constraints are best fulfilled by the solution with Teff=120000 K and log g =5.35.
  • we made an estimate of the external error as described in Napiwotzki (1999) and derived an error limit of 0.22 dex.
  • The position of PN G135.9+55.9 in the log g, Teff plane is shown in Figure 4.
investigating the periodicity of the radial variations
Investigating the Periodicity of the Radial Variations
  • We measured RVs for all CFHT spectra with temperature and gravity now fixed at the value Teff= 120000 K and log g =5.35 derived above.
  • used the Balmer lines Hγ to H9 and the two He ii lines for the RV determination.
  • The results: summarized in Table 1
  • The RV variation has a clearly periodic nature.
Figure 6 shows a ‘‘power spectrum’’ that was produced by fitting RV sinelike curves for a range of periods. The resulting fit quality, measured by χ2, is plotted
  • Apart from the formal χ2 values shown in Figure 6, simply consider the large RV shift between spectra 457 and 466: 470 km s-1 in just 2.3 hr!
  • ----the direct evidence that the nucleus of PN G135.9+55.9 is a close binary
  • We conclude that the orbital period cannot be longer than 4 hr, and a time series of spectra with much shorter exposure times is needed for an unambiguous period determination
the extinction of pn g135 9 55 9
  • From the FUSE data, we can estimate the extinction suffered by the object by fitting the flux distribution of PN G135.9+55.9 from the far UV to the near IR.
  • This is shown in Figure 8.
  • The slope of the observed UV spectrum (1000–1200 A) is consistent with the model, but shifted downward because of the interstellar extinction
  • we confirm that the reddening, and hence extinction, of PN G135.9+55.9 is indeed low. In the following, we adopt a value of E(B- V ) =0.04 mag

Fig. 8.—Flux distribution of PN G 135.9+55.9 from far UV to near-IR. The open squares are the observed fluxes. The dashed lines are continuum fits to the observed data. Fluxes corrected for interstellar extinction are represented by filled squares . The dot-dashed line is the expected contribution from the nebula . The open circles are the extinction-corrected data after subtraction of the nebular component. The central star model with 120,000 K and log g =5.35 is presented as a

solid line.

the distance and age of pn g135 9 55 9
  • The distance:
  • we can estimate the distance d to the object using the expression
  • FV is the model stellar flux in the V band in 108 ergs cm-2 s-1A-1
  • mV is the observed stellar magnitude corrected for nebular contribution and for extinction.
With different parameters, we can get different distances: d= 17.8, 22.6, 24.2, 20.6Kpc
  • In any case, our present estimates of the distance are consistent with PN G135.9+55.9 being located in the Galactic halo
The age
  • The expansion time, texp, of the planetary nebula is generally computed as

texp =Rout / vexp,

R out is the outer radius of the nebula and vexp is its expansion velocity

  • We adopt the nebular size determined by Richer et al. (2002) and an expansion velocity of 30 km s -1 as found by Richer et al. (2003).
  • Using this rough age estimate, the position of the object in the (Teff , texp) diagram (Fig. 9) can be compared with that of model stellar tracks of various masses
  • We find that the position of PN G135.9+55.9 is consistent with a mass of the ionizing star around 0.57 M
discussion the status of the stellar core of pn g135 9 55 9
  • The Ionizing Star
  • From the temprature and gravity we got above
  • From a comparison with evolutionary tracks of H-burning central stars of planetary nebulae ( Fig. 4) we estimated a mass of 0.88 Msun
  • This value is at variance with the expected properties of a Population II central star (the large height above the Galactic plane, the high radial velocity (RV 193 km s -1), and the very low metal abundance)
  • Because of the obvious Population II nature of PN G135.9+55.9, and the relatively large kinematic age of the nebula.---in the subsequent discussion we shall adopt a mass of 0.55 Msun
How to explain the apparent discrepancy between the observed and expected parameters of the ionizing star?
  • ----The progenitor almost certainly went through a common envelope phase, which drove it out of thermal equilibrium. In this case, the relations used to derive the stellar mass from log g and Teff do not apply.
  • ----the post-AGB star experienced a late thermal pulse (shell helium flash) after emerging from the common envelope phase.
The Companion Star
  • The short orbital period and large semiamplitude of the RV that we detect imply that the other component of the binary system should be a rather massive star.
  • The absence of a forest of RV variable emission lines in our spectra is another strong argument in favor of a compact companion, i.e., a white dwarf or a neutron star.
  • it is difficult to explain the formation of the PN G135.9+55.9 system with standard assumptions. On the other hand this means that PN G135.9+55.9 could provide an important test of our understanding of common envelope evolution.
If the companion is a compact star, the binary will merge in less than 1 Gyr. Our estimates from the mass function show that the companion should be rather massive and it is quite possible that the sum of the masses of this system exceeds the Chandrasekhar limit for white dwarfs (1.4 Msun)
young and very young stars in ngc 3372 the carina nebula
  • M. Tapia,1 M. Roth,2 R. Vazquez,3 and P. Persi
The Great Carina Nebula
  • Is a very important natural laboratory for studying the birth and evolution of the brightest and most massive stars in the Galaxy.
  • The stars are members of four distinct open clusters:
  • ----TR 15 to the farthest north,
  • ----Tr 14 and Tr 16 in the northern half of the nebula,
  • ----Cr 228, widely scattered in the southern region,
  • The present work presents a comprehensive photometric study of the northenr clusters, Tr 15, Tr 14 and Tr 16
  • Also Car I region

Fig. 1. R-band short-exposure mosaic of the northern Carina nebula region containing the clusters Tr 14 and Tr 16.

The circles mark the location and extension of the clusters Tr 14 and Tr 16 as derived from star counts.

  • All observations were carried out at Las Campanas Observatory in Chile using the 1.0 m (UBV RI and JHK), 2.5 m DuPont (JHK) and 6.5 m Clay/Magellan (deep JHK) telescopes
Assuming the clusters to be spherical, their centers and radii were determined by means of star counts. The results are shown in Table 1.
the dust extinction towards the stars in the Carina nebula is extremely variable,
  • ----both in terms of variations of optical depth (measured by AV )
  • ----and of variations of shape in the extinction law (measured by Rv = Av/E(B-V))
  • This is evidence of a very inhomogeneous intracluster dusty medium and of great diversity in dust particle size distributions across the nebula
Individual distance and extinction determinations were obtained to derive mean distances and reddening for the clusters (table 2).
There is large dispersion in both < d > and < AV > which are not caused by observational uncertainties. The data is compatible with the three clusters being a the same distance from the Sun, d = 2.7 kpc.
  • Star ages:
  • Analyses of the calibrated colourmagnitude diagrams indicate that
  • ----the stars in the dust-free cluster Tr 15 have ages between 3 and 40 million years
  • ----Tr 14 and Tr 16 are younger, with the older stars with ages around 6 million years and with new stars that are still being formed
For Tr 14 and its neighbouring radio HII
  • Near-infrared images reveal that the centrallycondensed cluster Tr 14 is partially embedded in adense cloud that lies behind the visible cluster.
  • Fig. 2 shows the distribution of thestars in Tr 14 as seen in H and K
  • Note the sudden decrease in star density occuring at the position of an ionization front expanding towards the densest part of the molecular cloud where, naturally, the extinction increases drastically
Extended radio emission from an HII region in the densest zones of the cloud suggest the presence of an embedded population of massive stars.
  • (by means of deep near-infrared imaging with the 6.5 m Clay/Magellan telescope.)
This paper presents the radial velocity structure of the molecular hydrogen distribution in the Ring Nebula (NGC 6720).
  • Observations of H2 v=1–0 S(1) ((2.122 um, ∆λ=0.02um) emission of the Ring Nebula were made in 2004 May with the 2.1 m telescope at the Observatorio Astrono´mico Nacional in San Pedro Ma´rtir, Baja California, Me´xico.
  • The measurements were carried out using the CAMILA infrared camera/spectrograph (Cruz-Gonza´lez et al. 1994) with a NICMOS3 array (256X256 ).
  • Velocity channel images of Ring Nebula H2 v=1–0 S(1) line emission obtained from the Fabry-Pe´rot data are shown in Figure 1
  • peak intensities are in the 0.6 to 19 km s-1 channel maps.
  • Shape: The emission in all the channel maps is contained in a ring, almost elliptical in shape, Whose size increases and decreases, with a maximum at the systemic velocity (+9.2 km s-1)
  • the H2 emission is not smooth, but rather highly clumped. Different knots appear at every velocity channel, and the size of a specific knot changes with velocity

Fig. 1.—Radial velocity channel maps of the 2.12 um H2 v=1–0 S(1) line emission in the Ring Nebula. LSR velocities in each channel are indicated in kilometers per second. The scale (in arcseconds) is given in the top left-hand panel. Maps show (in linear gray scale) the intensity of the H2 emission.. Crosses indicate the position of the central star. The noise level in all the images is constant, but the gray scale is varied in each

image to enhance the H2 emission.

A color composite of the H2 velocity structure of the Ring Nebula is shown in Figure 2.(created from mean value of velocity maps)
  • Blue: 10.4 and 0.6 km s-1
  • Green: 9.2 km s-1
  • Red: 19.0 and 28.8 km s-1
The green channel dominates most of the nebula, since it represent the channel with the highest emission intensity and the largest extension
  • A remarkable feature is the presence of the blueshifted bright knot in the northwest part of the ring.
  • H2 Morphology of the Ring Nebula
  • The molecular hydrogen emission of the Ring Nebula has a ringlike appearance that is approximately elliptical in shape and has a semimajor radius of 42”. and ellipticity of 0.70
  • Many authors claim that because of geometric considerations,
  • ringlike PNe displaying H2 emission are all bipolar nebulae that are viewed such that the polar axes are inclined with respect to the plane of the sky
  • However, this model does not
  • explain the large ellipticity of the central ring.
The channel maps in Figure 1 show that the molecular ring of the Ring Nebula does not present bipolar symmetry
  • To explain this phenomenon, we first assume a spherically symmetric, uniform density H2 shell expanding at a constant velocity.
  • If we approach a bipolar flow in which two of these spheres representing the red- and blueshifted flows, their velocity maps would consist of two sets like the one described above.
  • Since we only see a single progression, we conclude that ----the molecular ring of the Ring Nebula does not have a bipolar symmetry
If the axis of the H2 emission shell had a large inclination angle, then the blueshifted channel maps should be centered above the position of the central star, and the redshifted ring emission below the central star.
  • For the molecular hydrogen in the Ring Nebula, the centers of the blue- and redshifted emission rings are almost coincident (see Table 1).
  • ----the shell’s axis of symmetry is located along the line of sight.
Assuming an expanding velocity of the gas proportional to the distance to the central star, the position-velocity diagrams represent slices through the nebula perpendicular to the line of sight at fixed-distance intervals

Figure 3 shows a proposed model for the geometry of the H2 component of the Ring Nebula.

Thus, the best geometric model in accordance with the kinematic data is a hollow shell that is almost cylindrical in shape.
  • The axis of the cylinder lies along the line of sight.
  • The cross-section of the shell is elliptical
kinematics of the h2 knots
Kinematics of the H2 Knots
  • High-resolution images of the main ring of NGC 6720 show numerous condensations and filaments in the main nebula (Komiyama et al. 2000
  • It has been implied that the molecular gas is in some way shielded from the ionizing radiation because it is inside denser clumps
  • The origin of these molecular knots is uncertain.
  • A close inspection of the velocity channel images shows spots where the H2 emission is relatively intense.
  • This enhancement of the emission may be produced by dense concentrations of molecular hydrogen---- knots.
Figure 4 presents an H2 image of the Ring Nebula obtained by adding all the velocity channels.
  • The position offsets of the knots are shown in Table 2
  • The H2 v=1–0 S(1) line profiles of these knots are presented in Figure 5
  • Some of the profiles show the presence of several velocity components (knots A,and F), and others show a single-value peak (knots B, C, E,G, and H).
  • ---- This may suggest that some clumps are formed by one or more components that are differentiated by their kinematics.

Fig. 4.—H2 integrated image of the Ring Nebula obtained from adding all the velocity channel maps. Labels indicate the position of the molecular knots detected in the velocity channels (see Table 2).


Fig. 5.—Profiles of the H2 v=1–0 S(1) line emission obtained at the positions shown in Fig. 4 and Table 2, with a spectral resolution of 24 km s-1. The line profiles are plotted on arbitrary scales that are not an indication of the relative intensities of the different lines.

  • it has been shown that the knots have distinct kinetic characteristics.
  • The different radial velocities observed in the knots prove that they are located all around the ellipsoidal shell, and not just in the equator of the main nebula
co and h2 kinematics
CO and H2 Kinematics
  • The CO kinematics of the Ring Nebula have been investigated by Bachiller et al. (1989).
  • Because their spectral resolution is higher (1.3 km s-1 at 230 GHz); they complement the observations presented in this work.
  • the discussion is based on their results for CO J=2-1 lines.
  • The CO emission extends on a ring that resembles the H2 and optical images
  • The CO and H2 images are clumpy, with the emission peaks of the clumps located at roughly the same position.
  • The spatial distribution of the CO emission is roughly compared to the H2 emission velocity integrated map in Figure 4.
  • comparing the CO profiles J = 2–1 at offset positions (-40”, -20”) and (0”, 30”) to the H2 knot counterparts C and A,
  • ----find that The CO and H2 emission profiles at each of these positions have similar overall shapes and velocity widths.
  • The comparison of velocity maps for both emissions show kinematic similarities that imply that molecular hydrogen and CO are well mixed throughout the nebula.