structure of the orion nebula

Structure of the Orion NebulaAuthor(s): C. R. O’DellSource: Publications of the Astronomical Society of the Pacific, Vol. 113, No. 779 (January2001), pp. 29-40Published by: The University of Chicago Press on behalf of the Astronomical Society of the PacificStable URL: http://www.jstor.org/stable/10.1086/317982 .

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29

Publications of the Astronomical Society of the Pacific, 113:29–40, 2001 Januaryq 2001. The Astronomical Society of the Pacific. All rights reserved. Printed in U.S.A.

Structure of the Orion Nebula

C. R. O’Dell

Department of Physics and Astronomy, Box 1807-B, Vanderbilt University, Nashville, TN 37235; [email protected]

Received 2000 September 28; accepted 2000 October 12

ABSTRACT. The nature of the structure of the Orion Nebula has become clearer but more complex. Quantitativeapplication of simple photoionization theory has allowed construction of a three-dimensional model of the mainionization front, while high-resolution study of the flow of [O i], [O ii], [O iii], and [S iii] has allowed determinationof how this material moves. Material in a “foreground lid” of H i is seen in absorption lines in the 21 cmcontinuum and in optical spectra of Na i and Ca ii. There remain many unsolved and possibly basic questions,among them the source of nonthermal broadening of all lines, which carries as much energy as the thermalcontent of the gas. We have also found that varying amounts of emission-line light scattered by the dust particlesimmediately behind the main ionization front introduces a nonphotometric scatter of up to 25% in spectralintensities.

1. INTRODUCTION

Our understanding of the physical structure of the OrionNebula has continued to evolve and become more detailed thanat the time it was described in the Symposium on the OrionNebula to Honor Henry Draper by M. Peimbert (1982). Bythat time the current basic model, where the nebula that wesee as M42 is a thin photoionized layer on the surface of thegiant molecular cloud OMC-1, had been widely accepted. Themodel was advocated at about the same time by several authors(Zuckerman 1973; Balick, Gammon, & Hjellming 1974; Bal-ick, Gull, & Smith 1980; Meaburn 1975; Pankonin et al. 1979);Zuckerman and Balick et al. (1974) are usually given creditfor the defining argument in favor of the model, that one ob-serves a gradient of progressively more blueshifted emissionlines, with the highest ionization lines being the most blue-shifted. In fact, a slab model for the emission goes back muchfarther. It was originally invoked by Munch (1958) to explainthe statistical properties of the fine-scale motions across theface of the nebula and also by Wurm (1961) to explain themonochromatic emission-line and broadband scattered light im-ages that he was obtaining (Wurm & Rosino 1959, 1965). TheWurm article appeared only in German and apparently was notwidely read, and Munch & Wilson (1962) later advocated anonslab model, so that the question of the three-dimensionalstructure remained unclear up until the middle 1970s. Once theparadigm shift occurred, many known features fell into place.Progress since the 1981 meeting has been considerable, largelydue to the increase in the number of investigations of a quan-titative nature, using CCD detectors for both imaging and high-resolution spectroscopy. These studies show that M42 has amyriad of features not anticipated two decades ago.

In the succeeding sections I will discuss what we know of

the main ionization zone (§ 2), the nature of the lid that coversthe front (§ 3), the bright bar (§ 4), velocities in the nebulaand its associated stars (§ 5), and the role of scattered light(§ 6); I then summarize with a discussion of unresolved prob-lems (§ 7). Questions of the exact physical conditions, such aselectron temperature, are discussed in a companion article byFerland (2001), who presents a description of the closely relatedsubject of the atomic abundances.

2. THE MAIN IONIZATION ZONE

The paradigm shift that occurred nearly 30 years ago wasthe recognition that M42 was a thin blister of photoionized gason the surface of OMC-1. Such a structure would be dynamic,in the sense that the overpressure situation produced by heatingof the gas would cause material to be lost by almost uncon-strained expansion. This process is often labeled as a cham-pagne phase flow, although that is not rigorously correct, sincethe champagne phase flow more exactly describes the emptyingof a previously enclosed H ii region when its constrainingneutral boundary is breached. However, the Tenorio-Tagle(1979) models do capture the important element, that an ionizedlayer at the surface of a giant molecular cloud is a nearly staticphenomenon, rather than a fixed feature. The molecular cloudhas a large reserve of material, in this case of about 103 M,.New molecular cloud material is added to the ionized gas assome is lost through expansion. There is theoretical andobservational evidence (Tielens & Hollenbach 1985; Escalante,Sternberg, & Dalgarno 1991) that a shocked, highly com-pressed zone (the photon-dominated region, or PDR) existsimmediately behind the ionization front (IF), meaning that thereis so much material that the specific location of the exact IFwill change only slowly. However, the flow of gas away from



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TABLE 1The Ionization Front behind v1 Ori C

Zone PDR IF Low Ionization Medium Ionization

Key ion . . . . . . . . . . . . . . H0 H1 He0 He1

Markers . . . . . . . . . . . . . . CO, C ii [O i], [S ii] [O ii], [N ii] [O iii], H ii, He i, [Cl iii]V,

a (km s21) . . . . . . . . . 28 25.5 18.8 5 1.5 17.9 5 1.3Densityb (cm23) . . . . . . 105 ≥6000 7000 4000Depthc (pc) . . . . . . . . . . . ? 1024 2 # 1023 0.06

a Velocities are from Goudis 1982; O’Dell & Wen 1992; and Hu 1996.b Densities are from Tielens & Hollenbach 1985; Escalante et al. 1991; Pogge et al. 1992; Jones

1992; and Walter 1993.c IF and low-ionization depths are from O’Dell 1994.

the IF may not be absolutely unconstrained, since the radiationpressure from the stars of the Trapezium approximately equalsthe pressure at the base of the ionization front, which meansthat if the material becomes optically thick to the stellar softradiation, then ionized material may expand more slowly at thesubstellar point. We now have the radial velocity material totest whether an exact theoretical model with and without partialradiation pressure confinement best fits these observations.What is lacking is the realistic theoretical calculation for theconditions in M42. The expectation is that material would beaccelerated away from the IF, i.e., become more blueshiftedwith respect to the observer, and that the density should de-crease as this acceleration occurs. Independent of the detailsof this process, it is important to understand that this spectacularobject is produced with only about 2 M, of material (Wilsonet al. 1997).

There should be an ionization stratification that occurs withinthe ionized level. Fortunately this is easy to predict since thereis only one important source of ionizing photons and the opacityof UV photons is primarily determined by only hydrogen andhelium. The spectral type of v1 Ori C, the brightest and hotteststar in M42, is O7 (Conti & Alschuler 1971), which meansthat its photoionizing luminosity (above 13.6 eV) is about 3–4times greater than the closest competitor v2 Ori A, which isspectral type O9 V (Conti & Alschuler 1971) and lies awayfrom the Trapezium. As pointed out in Osterbrock’s (1989)text, the range of absorption coefficients of various atoms andions varies little as compared with the abundances of the el-ements. This means that radiative transfer in the region im-mediately above 13.6 eV will be dominated by hydrogen ab-sorption, that above 24.6 eV by He i, and that above 54.4 eVby He ii. This means that the heavy elements such as carbon,nitrogen, and oxygen, which produce many observable lines,should be found in zones whose primary structure is determinedby the ionic state of H and He. The ionization front itself willgive rise to [O i] and [S ii] emission, since they require boththe presence of these low-ionization states and electrons ofseveral eV to cause the collisional excitation of the forbiddenlines. The next zone out from the IF will have H1 but neutralhelium (He0) and produce emission from [N ii] and [O ii].

Outside of that zone the helium will exist as He1, and itsemission in recombination will be accompanied by collisionallyexcited [O iii] and [Cl iii]. Higher states of ionization are notimportant because of the temperature of v1 Ori C. This simplestratification of ionization states agrees well with our obser-vations, with Table 1 giving the range of velocities and electrondensities that are observed, which agrees well with an an emit-ting layer of stratified ionization, with the material farthest fromthe IF having the lowest densities and the greatest blueshifts.

For several decades the discrepancy between the high den-sities measured by indicators such as the [O ii] l3727 doubletand the surface brightness [S(Ha)] in Ha was resolved byintroducing a “filling factor.” The need for this concept dis-appears with the recognition that the emitting layer is thin ascompared with the overall size of the H ii region and that, infact, one can calculate the thickness of the emitting layer. Bald-win et al. (1991) first showed how their calibrated slit spectrasurface brightnesses could be used to calculate a formal thick-ness for the emitting layer, when the local density was known.This method was extended by Wen & O’Dell (1995) acrossthe brightest parts of the nebula, and they found that the layerthickness was about 0.13 pc under the assumption that thedensity is constant. If the density is assumed to fall exponen-tially, the e21 scale height is half of this value. Since the em-issivity of all of the observed lines scales as the square of thedensity, this means that the e21 scale height of the observedemission is only pc; this corresponds to an an-0.13/4 p 0.03gular distance of 140, which is small when compared with thedimension of the nebula (it is 1350 between v1 Ori C andv2 Ori A) if the distance to the nebula is the same as that ofthe associated stars, 430 pc (Warren & Hesser 1977).

Ferland (2001) also showed in the Baldwin et al. (1991)paper that one could also calculate the separation distancebetween the photoionizing star and the IF. In the case ofphotoionization equilibrium by a thin layer at a distance rfrom a star of ionizing luminosity Q (photons s21), the totalnumber of recombinations in a column of material must beexactly equal to the photon flux ( ). Since a fixed frac-2Q/4prtion of the recombinations will produce an observable Ha

photon, then the extinction-corrected Ha surface brightness



STRUCTURE OF ORION NEBULA 31

2001 PASP, 113:29–40

in photons s21 sr21 will be

effa Q 1HaS(Ha) p , (1)2 2a 16p rB

where and aB are the hydrogen recombination coefficientseffaHa

defined and given values by Osterbrock (1989). Ferland’sformulation is rigorously valid only for the substellar point,where the photons enter along the same line of sight as theobserver’s view; the basic process was expanded to cover theentire inner region by Wen & O’Dell (1995), who made de-tailed calculations of the effect of the incident and observeddirections when they differed. This required knowledge ofthe density distribution across the nebula, for which they util-ized the densities determined from the [S ii] doublet ratiosby Pogge, Owen, & Atwood (1992). They found that theseparation at the substellar point was about 0.3 pc and thatthe nebula was a highly irregular concave surface, which isunusually high to the southwest of v1 Ori C, where the nebulais brightest, and curves up abruptly along the northeast-south-west bright bar that runs just north of v2 Ori A. Their substellarseparation is confirmed by an independent calculation fromthe surface brightness of the nebula in scattered starlight(O’Dell 1994).

The bottom-line result of these calculations is that the nebulais really quite a low-mass object (2 M,), whose primary emis-sion occurs from a glowing layer that is thin (0.03 pc) ascompared with the minimum stellar distance (0.3 pc) and thecharacteristic dimension of the object in the plane of the sky(0.6 pc, which is about twice the v1 Ori C–v2 Ori A distance).

3. THE FOREGROUND LID

One of the most significant investigations concerning thestructure of the M42 region since the Henry Draper Symposiumwas that of van der Werf & Goss (1989). In this study theyused the 21 cm thermal continuum emission from M42 as abackground source and searched for H i absorption lines. Theregion covered was centered on the Trapezium stars and ex-tended to a diameter of about 109. Superb velocity resolutionwas employed (0.64 km s21), and the spatial resolution (180)was sufficient to resolve most of the features detected. Clumpedneutral hydrogen absorption was detected in front of much ofthe nebula, generally showing an increase of column densityN(H i) toward the northeast, where the optically dominant DarkBay feature is seen to the east of the Trapezium. A follow-upinvestigation was made by O’Dell et al. (1992) to determineif the H i column density correlates with the amount of ex-tinction. Extinction was determined by comparing the Ha/Hb

ratio point by point from calibrated images (Hester et al. 1991)of the nebula. It was found that there was an excellent corre-lation of N(H i) with extinction, which indicates that most of

the extinction suffered by the nebula and its nearby clusterarises in the material producing the 21 cm absorption.

The location of this material is hard to pinpoint. However,it probably lies quite close to the nebula as the velocities ofthe material are close to but always blueshifted with respect tothe OMC-1 velocity. None of the velocities have the high blue-shifts associated with the outermost parts of the entire OrionComplex which are commonly referred to as Orion’s Cloak(Cowie 1982), so that the new H i material is probably bestcalled the “foreground lid” to M42. “Veil” might be a moreaccurate description, since the visual optical depth is about 2,diminishing to the southwest and becoming much larger in theDark Bay, which is a continuous part of the Foreground Lid.It is difficult to place the Lid exactly. The S(Ha) versus stellardistance equation should also apply to the IF that lies on theside of the Lid closest to v1 Ori C, and a considerable efforthas been made to detect emission from the lid. Unfortunately,the thermal width of the Ha lines is about 15 km s21, whichis larger than the velocity difference of the Lid from the mainIF, with the result that emission from the inside of the Lid hasnot been detected. There is an extended region of [S ii] emissionwhich may mark that IF, but one cannot locate the positionfrom its intensity. Recent work on Herbig-Haro type shocks inthe vicinity of the Trapezium cluster indicates that a numberof shocks, such as HH 203–204 and HH 269, are featuresmarking where outflow from cluster stars are colliding withthe Lid (O’Dell et al. 1997b), and the placement of these objectsargues that the Lid separation is approximately the size of theoptical nebula. This means that the Lid lies about 0.6 pc infront of v1 Ori C. The Lid seems to be composed of threesystems, each distinct in its characteristic velocity even thoughhighly irregular in N(H i). In the terminology of van der Werf& Goss, the systems and their heliocentric velocities are A(24 km s21), B (21 km s21), and C (16 km s21). These threesystems are also evident in a study of the interstellar absorptionlines of Ca ii in the four Trapezium stars plus v2 Ori A (O’Dellet al. 1993) and of Na i in v1 Ori C and v2 Ori A (Hobbs 1978).The spectra of these heavy elements show multiple velocitycomponents, with coinciding Ca ii components at 22.0, 19.8,and 15.3 km s21 and Na i components at 23.0, 18.3, and 14.9km s21. All of the absorption-line data are presented in Table2, and the velocity systems identified with the Lid are desig-nated there as V7 and V8. The additional noncoinciding com-ponents are discussed in § 7.6. Even though neither ion is themost populous species of the atom, they produced measurablelines because one can observe transitions from the highly pop-ulated ground state. Ca ii is produced by 6.1 eV photons andis in turn photoionized by 11.9 eV photons, which means thatCa iii is the most populous species of calcium, since photonsof less than 13.6 eV will penetrate into the neutral hydrogenclouds. A similar situation applies for Na i, which is photoion-ized by 5.1 eV photons. It should be possible to determine theapproximate displacements from v1 Ori C from the relative



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TABLE 2Absorption Lines in the Trapezium Stars and v2 Ori A

Region Ion V1 V2 V3 V4 V5 V6 V7 V8 V9 V10

v1 Ori A . . . . . . He i … … … … 3a … … … … …Ca ii … … … … 2 10b 17 23b 31b 40

v1 Ori B . . . . . . He i … … … … 1a … … … … …Ca ii … … … … 23 6b 17b 23 31b 40

v1 Ori C . . . . . . He i 260 230 … 211 2a … … … … …Ca ii … … 217 … … 7b 14 21b 31b 39Na i … … 217 … … 6b 17b 22b 31b 38

v1 Ori D . . . . . . He i … … … … 3a … … 21 … …Ca ii … … … … … 7b … 22b 31b 40

v2 Ori A . . . . . . He i … … … … 23, 15 … … … … …Ca ii … … … … 1b … 14b 20, 23b 31b 39Na i … … 215 … 1b 11b 15b 24b 32 …

All . . . . . . . . . H i … … … … … … 16 21, 24 … …

a He i velocities with equivalent widths greater than 100 mA.b Ca ii and Na i velocities with equivalent widths greater than 15 mA.

abundances of Na i and Ca ii, which have been measuredaccurately.

Van der Werf & Goss (1990) also report the presence of afew rapidly moving H i clouds, characteristically blueshiftedabout 15 km s21 with respect to the Lid systems. They explainthese as having been accelerated by the rocket effect ofphotoablation, where clouds much more massive would nothave been accelerated enough to be displaced and detected andclouds less massive would have been destroyed by photoion-ization. Unfortunately, no optical counterparts to these cloudshave been seen. This could be used as an argument against therocket acceleration interpretation; however, the ionizationboundaries producing the push would be weak as comparedwith the background nebular emission of the main IF.

4. THE BRIGHT BAR

An IF will typically produce a shocked high-density regionon its neutral side (the PDR). This region will have a gradientof conditions basically determined by the extinction in theultraviolet continuum on the low-energy side of 13.6 eV, thecutoff imposed by H i. As nicely described by Tielens et al.(1993), this leads one to expect a series of layers. The closestto the IF will be C ii as a result of seeing photons in the11.3–13.6 eV range (Wyrowski et al. 1997), with the next layerbeing of H2, and the deeper layer being optimally traced byCO emission. They also point out that the polycyclic aromatichydrocarbon molecules, with a strong emission feature at3.3 mm, will also occur in the region immediately behind theIF. All of these spectral features are observed when lookingat M42 in the radio and infrared regions, but the interpretationof the data is usually not clear because one is looking along aline of sight passing about perpendicular to the IF. Fortunately,this is not the case for the Bright Bar, a region of enhancedHa surface brightness that runs about 69 from the northeast tothe southwest and passes about 300 north of v2 Ori A. Not onlyis this remarkable linear feature brighter, but it also shows an

enhanced [N ii]/Ha ratio, generally being greater on the sidefacing v1 Ori C. This led to the early recognition that the BrightBar is caused by the main IF being tipped almost along theobserver’s line of sight (Dopita, Dyson, & Meaburn 1974;Balick et al. 1974). This position has been taken as a point ofdeparture by each study since then. Wen & O’Dell’s (1995)three-dimensional model of the nebula showed that this viewof the Bright Bar is basically correct; however, their calcula-tions indicate that the IF is not exactly perpendicular, althoughit must be admitted that their method tends to smooth the nebulaand reduce the gradients in the surface. A detailed study ofHubble Space Telescope WFPC2 images along the bar (O’Dell& Wong 1996) shows that there are a number of peculiar fea-tures, which include crenellated structures (O’Dell & Wen1994) on the side away from v1 Ori C (which should not beobserved if we are looking at an IF viewed exactly along itsedge) and even regions where the [O iii] emission is enhancedon the side away from v1 Ori C. Although the first-order struc-ture is caused by our viewing an IF almost edge-on, there areobviously many questions to be resolved.

5. STRUCTURAL FEATURES DERIVED FROMVELOCITY OBSERVATIONS

The fact that the emission-line radiation is packed into anarrow wavelength interval has allowed the application of veryhigh spectral resolution in some studies of the nebula, and thebrightness of the Trapzeium stars and v2 Ori A has also allowedsimilar high resolutions for investigation of absorption linesformed in material lying between the observer and those stars.The FWHM broadening of an ion of atomic weight M in a gasof electron temperature 8000 K will be FWHM p 19.2 #

km s21 and will increase as the square root of the tem-21/2Mperature. This means that resolutions of a few km s21 are ashigh as necessary for studying the hot nebular gas, althoughbetter resolution will benefit absorption studies of a cooler gas.

The ultimate data set of velocity information would be cov-




2001 PASP, 113:29–40

erage of every major ion at a spectral resolution that wellresolves the line profiles and at a spatial resolution matchingthat of the best images. We are well short of that goal but areslowly making progress. A series of papers by the author andhis students has provided adequate velocity resolution studiesin [O iii] (O’Dell & Castaneda 1987; Castaneda 1988),[O ii] (Jones 1992), [O i] (O’Dell & Wen 1992), [S iii] (Wen& O’Dell 1993), and [S ii] (Hu 1996) using slit spectroscopy.Although multiple slit settings were employed, the spectracover only a small, but probably representative, sample of thecentral 59 diameter region of the nebula. Radio observationsof H ii have covered the entire nebula at superb velocity butlow spatial resolution (Wilson et al. 1997). The central 109region has been covered continuously at 50 km s21 resolutionin [O iii] and [S ii] using a Fabry-Perot imaging system (O’Dellet al. 1997a). This latter study is quite adequate for identifyingvery high velocity features, which abound as a result of out-flows from many pre–main-sequence stars. However, these in-teresting features do not play an important role in the dynamicsof the nebula as a whole and fall outside the scope of thepresent paper.

Only the brightest stars can be used for study of absorptionlines caused by intervening gas. The early studies of v1 Ori Cand v2 Ori A (Munch & Wilson 1962; Munch 1985) haverecently been supplemented by studies of all four members ofthe Trapezium (Oudmaijer et al. 1997) and the Trapezium mem-bers plus v2 Ori A (O’Dell et al. 1993).

5.1. The Large-Scale Velocities of M42

The entire nebula has been mapped in the H64a line at1.4 km s21 and 420 by Wilson et al. (1997) using the 100 mtelescope of the Max-Planck-Institut fur Radioastronomie. Thecentral 39 diameter region was also studied using the BalmerH12 optical line by Jones (1992) at a resolution of 200 withvery similar results. Since the H ii emission comes from allionized zones, it is the least discriminating ion for study of theacceleration of material away from the IF but the best guideto the overall flow of gas. The radio maps have a particularadvantage in studying the flow in the eastern part of the nebulawhere the Dark Bay extinguishes much of the optical radiation,especially so as the nebula also becomes of lower intrinsicsurface brightness there (Yusef-Zadeh 1990). The H ii datashow that there is a central most positive velocity of about20 km s21 centered about 300 east-southeast of v2 Ori A.Throughout this paper all velocities will be heliocentric, whichare 18.3 km s21 more positive than the local standard of restvelocities favored by radio astronomers and students of Galacticstructure. There is a global trend for the velocities to becomeabout 6 km s21 more blueshifted as the radial distance increasesfrom v1 Ori C out to a distance of about 29.5. In addition, thereare two nonradially symmetric trends. About 39 to the south-west the flow is an extra 2 km s21 more blueshifted. About 39due west of v1 Ori C, well behind the Dark Bay, the velocity

abruptly shifts back to the central values and perhaps becomesslightly more positive (see Fig. 1).

The optical studies are limited in the area of coverage bythe decreasing surface brightness away from v1 Ori C and alsothe extinction due to the Lid material. The central 29.5 diameterregion shows an almost constant velocity at 25.5 km s21 in[O i] (O’Dell & Wen 1992). The [O ii] emission has a centralvalue of about 19 km s21 and becomes several km s21 bluerto the southwest of v1 Ori C. [O iii] emission is about 22 kms21 in the central region, generally becoming a few km s21

more blueshifted at increasing radial distance with an asym-metry of even greater blueshift to the southwest.

Within the paradigm that all of the neutral gas has the ve-locity of OMC-1 (28 km s21) and that there is a local flow ofgas perpendicular to and away from the local IF, then one canuse these velocities to discuss the three-dimensional structureof M42. The radially symmetric increasing blueshift with dis-tance from v1 Ori C is the opposite to what one would expectfrom a concave surface (which is indicated by the three-dimensional model of Wen & O’Dell). It is as if the rate ofacceleration of material away from the IF is much lower in theregions closest to v1 Ori C, which probably is the case, sincethe radiation pressure from the stellar continuum will be highestthere and the gas will be held back. The asymmetric bluewardflow to the southwest has been interpreted by O’Dell & Wen(1994) and Wilson et al. (1997) as an indication that this iswhere the ionized material is escaping the confinement of theLid, which is consistent with the fact that the Lid becomesvery thin or perhaps even vanishes in that region. The abruptincrease of H ii velocities 39 west of v1 Ori C was interpretedby Wilson et al. as being due to the IF curving abruptly towardthe observer, thus removing the expansion component of ve-locity. This interpretation is consistent with the 21 cm contin-uum map of van der Werf & Goss (1989), at 220 resolution,which shows the surface brightness diminishing less with dis-tance. Since the radio surface brightness contours behind theDark Bay region (van der Werf & Goss 1989; Yusef-Zadeh1990) decrease asymmetrically fast, this indicates that the IFis less concave than in the optically visible region. The ob-servations of Wilson et al. may be showing us the region wherethe IF curves around toward the observer and joins up withthe Lid.

5.2. Turbulence in the Main IF

The availability of the slit spectra of M42 has made it pos-sible to investigate the statistical properties of the main IFvelocities and to compare these with the expectation of tur-bulence theory. Each long-slit spectrum could be divided intomultiple spectra, each representing a spatial sample of about

, thus providing an adequate sample of several hundred′′ ′′2 # 5accurate points. The theoretical framework has existed fornearly half a century. Von Horner (1951) solved the problemof relating observable radial velocities to the intrinsic distri-



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Fig. 1.—This is a depiction of the center of the Orion Nebula Cluster and M42 as viewed from the southwest. The escarpment which produces the Bright Barfeature is on the far left, and the brightest region, which lies over the Orion S molecular outflow source, is in the low middle. The relative positions of the starsand proplyds have been determined from the assumption of a symmetric distribution in three dimensions as determined by Hillenbrand & Hartmann (1998) exceptwhere there is specific evidence for the position. The shocks associated with outflows from low-mass stars are shown in positions developed in O’Dell & Bally(2000) and Bally, O’Dell, & McCaughrean (2000). The distance to the inner side of the Foreground Lid (horizontal bright feature above the brightest star, v1 OriC) is probably the most uncertain dimension in this depiction. This image was rendered by the San Diego Supercomputer Center under contract with the AmericanMuseum of Natural History as part of the Three-D Galaxy project of the new Hayden Planetarium.

bution of velocities, where it was assumed that the relativevelocity ( ) of two points separated by a distance (r) will bev

. For the common case of Kolmogorov turbulence,2 nv ∝ r. Von Horner’s model was that of a thin slab nebula,2n p 3

and he determined the properties of the radial velocityderived along various lines of sight. The cleanestV(f)

method of analysis is through the statistical property calledthe structure function [ ], which is defined asB(f)

, where f refers to the angular sep-′ ′′ 2B(f) p AFV(f ) 2 V(f )F Saration of two samples located at positions f9 and f0 and

is calculated for all combinations of velocity samples. AtB(f)small values of f, ; at large values, .n21 nB(f) ∝ f B(f) ∝ f

The transition to “large” occurs at the angular distance that isequivalent to the thickness of the emitting layer.

The value of was calculated for the radial velocities ofB(f)each of the ions ([O i], [O ii], [O iii], and [S ii]) that wasstudied with the slit spectrograph. Three of the ions showed

the theoretically predicted steep power law at small f-valuesand a transition to flatter power laws at large f-values. [O ii]and [S iii], which should arise from similar levels of ionization,were very similar, being slopes near unity for small f-valuesand essentially constant for large f-values, with the transitionoccurring at about 220. [O iii] made a transition from a powerlaw of 0.8 to one of 0.3 at about 150. It is remarkable that thetransitions occurred at about the same f, since the [O ii] layershould be very thin as compared with the [O iii] layer, as shownin Table 1. Moreover, the well-defined large f power laws weremuch less than the Kolmogorov theory value of . One can2

3

summarize these results by saying that these ions show resultsthat resemble, but do not agree with, the predictions of vonHorner’s models. This result is different from the older analysisof Munch (1958), who applied a similar technique to a set ofphotographic M42 spectra (Wilson et al. 1959).

The one ion whose structure function agrees well with the




2001 PASP, 113:29–40

prediction of the model is [O i], which shows a constant powerlaw of over the entire range of f that was observed0.7 5 0.1(60–1800). The power law agrees with Kolmogorov theory, andthe lack of a break in the slope is consistent with even thesmallest values of f being much larger than the thin [O i]emitting layer (Table 1).

There is, however, another discrepancy with theory whichapplies to all of the ions observed. According to the turbulencemodel there will be a turbulent broadening of the emission linealong each line of sight, effectively representing random mo-tions along a line. The magnitude of that broadening canbe estimated from the absolute value of , and in eachB(f)case the turbulent broadening should be a small fraction of1 km s21. Von Horner and Munch found the same result forthe early [O iii] observations. In fact, each of the lines studiedshows a nonthermal broadening of about 9 km s21. By non-thermal, I mean that after one quadratically subtracts the ex-pected thermal broadening, there always remains a componentof broadening with km s21. A very similar resultFWHM . 9was found by Wilson et al. (1997) for H i radio emission.

Such an extra line broadening may be very significant, be-cause it means that along each line of sight as much or moreenergy is being carried by this nonthermal turbulence as iscontained in the Maxwellian random motion of the individualatoms. However, it is not obvious that this is important indetermining the energy balance of the nebula. We know thatthe thermal gas loses its energy in a characteristic time of about

years, where is the electron density (cm23), which100,000/N Ne e

means a timescale of about 10 years for the densest part ofM42. Without knowing the mechanism producing the extrabroadening, we cannot derive the power necessary to sustainit. O’Dell & Wen (1992) argue that it is due to fine-scaleirregularities in the shape of the main IF, an interpretation dis-cussed in another context in § 7.3.

5.3. The He i Absorption-Line System

The presence of a metastable level in He i gives rise to alarge optical depth in some He i emission lines formed nearthe main IF and also the potential for formation of absorptionlines in the stellar continua. The lowest lying triplet state(2 3S) is populated by He1 recombinations and depopulated bytwo-photon emission to the ground-state singlet, collisions tothe 11 P state, and photoionizations by Lya photons of hydro-gen. Very early observations of v1 Ori C revealed the He il3889 line in absorption, which has been followed by manysubsequent observations (Munch & Wilson 1962; O’Dell et al.1993; Oudmaijer et al. 1997). The 3889 A line is formed bythe 2 3S–3 3P transition and is instrinsically weaker than themore difficult to observe the 10830 A line of the 2 3S–2 3Ptransition. However, improved near-infrared detectors have ledto the measurement of that line too (Vaughan 1968; Munch1985; Oudmaijer et al. 1997). Our primary source of infor-mation lies with the 3889 A line, although any observation of

the central stars is made against the He i emission lines pro-duced by the nebula. The much greater brightness of v1 Ori Cat 3889 A more than compensates for its lower oscillatorstrength, although Munch & Pitz (1981) have used a choppingtechnique that compensates for much of the nebular emissionat 10830 A. This contrast effect means that the determinationof the absorption spectrum will be less well determined for thefainter stars in the Trapezium.

The two modern CCD studies (O’Dell et al. 1993; Oudmaijeret al. 1997) have produced similar results for the four Trape-zium stars, while only the earlier study included v2 Ori A.O’Dell et al. used a velocity resolution of 3.3 km s21, whilethat of Oudmaijer et al. was ∼7 km s21. Table 2 gives theaverage results of these studies for the Trapezium stars andO’Dell et al.’s values for v2 Ori A. The component reportedby Oudmaijer et al. at 260 km s21 in v2 Ori A was also seenin the O’Dell et al. study, although it did not reach their cred-ibility threshold for publication. It is certainly there, as it ap-pears in these two very independent studies. The Oudmaijer etal. values for this line are given in Table 2.

The primary feature is at 3 km s21 which is common to allof the Trapezium stars. The 3889 A line in v2 Ori A has ab-sorption features close to this velocity, but very clearly thev2 Ori A line requires two components to fit the line profile.

O’Dell et al.’s spectra also allowed measurement of thenearby [Ca ii] line, and Hobbs (1978) measured [Na i] inv1 Ori C and v2 Ori A. O’Dell et al. found a strong [Ca ii]component at 7.5 km s21 in the Trapezium stars and anotherat 0.5 km s21 in v2 Ori A. These are also included in Table 2,where the various systems are grouped by similar velocities.

5.4. Where Are the Absorption Lines Formed?

Only the two velocity systems (7 and 8) associated with theforeground lid seem to have an established region of formation.The origin of the other velocity systems is quite uncertain withthe exception of the 21 km s21 component of He i, which couldarise from absorption in the main IF. This would mean thatv1 Ori D would have to lie closer to the main IF than the otherTrapezium stars, an interpretation that is consistent with thefact that the Ney-Allen infrared source is centered on this star.

There is no He i emission-line counterpart of any of theHe i absorption systems, which is not surprising since it takesvery few atoms in the 2 3S state to produce the observed l3889line. This argument is most clearly given in Baldwin et al.(1991) and repeated in a slightly different form by Munch &Wilson (1962) and Oudmaijer et al. There are emission-linecomponents of [O ii] near the Trapezium at 3.1 km s21 (Jones1992) and of [O iii] at 2.8 km s21 near v2 Ori A (Castaneda1988), so perhaps those are arising from the same sources.

The most difficult thing to decipher in interpretation of theabsorption lines is that system 5 seems to produce both thehigh-ionization He i lines and the low-ionization Na i andCa ii lines. If indeed all of these velocity features have the



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same region of origin, it will be a parcel of material if they liewithin the boundary of the main IF (on the back) and the Lid(on the front) unless they are short-lived. A relative velocitywith respect to OMC-1 and the main IF of 25 km s21 meansthat 1 pc is traversed in about 40,000 years, a time short com-pared with the yr age of the cluster. There is6(0.5–1) # 10indeed a wealth of fine-scale features in the nebula that movewithin this velocity range. These are the shocked material pro-duced by outflows from low-mass pre–main-sequence stars,some of which have previously been detected as Herbig-Haroobjects, while most of them have been revealed through thenew HST images (O’Dell et al. 1997b) and Fabry-Perot dis-crimination of their high velocities (O’Dell et al. 1997a). Theseshort-lived features have sizes from a few arcseconds to manytens of arcseconds, which means that they could cover the entireTrapezium region and produce multiple velocity systems. Theidea of producing observable absorption lines of Na i andCa ii in shocks within an H ii region seems to have not beenconsidered, so that it is difficult to say if this interpretation isplausible.

5.5. The Relative Velocity of Various Systems

Membership and association of the stars in Orion are usuallydetermined from proper-motion data. These are not availablefor the nebular components of M42 and are notoriously in-accurate for the Trapezium stars because of their brightness, acondition that is not changed by the new Hipparcos data be-cause the four bright stars of the Trapezium are too close to-gether. One can, however, obtain some useful information fromthe radial velocities. The radial velocity of the molecular com-ponents of OMC-1 is 28 km s21 (as determined from the com-pilation in Table 3.3.VII of Goudis 1982). The nebular valuesare as blue as 25.5 km s21 at the IF, with the main body atabout 18 km s21. The average of v1 Ori A, v1 Ori B, andv1 Ori D is km s21 if one uses the values from Abt,24 5 3Wang, & Cardona (1991). Both v1 Ori A and v1Ori B arespectroscopic binaries, where one must use their systemic ve-locities, while v1 Ori D is a broad-line B0.5 V star. The low-mass stars of the cluster give an average radial velocity of

km s21 (R. Mathieu & L. Marschall 1998, private com-27 5 3munication). All of these numbers are consistent with the clus-ter having no more than a few km s21 relative velocity withrespect to OMC-1 and the fact that the nebula’s blueshift is aresult of expansion of the ionized gas away from the surfaceof OMC-1.

5.6. The Riddle of the Radial Velocity of v1 Ori C

The radial velocity of v1 Ori C is extremely important inunderstanding what has happened, what is happening, and whatwill happen in M42. The recombination time for photoionizedhydrogen is about 100,000/N years (where is the electronNe

density in units of cm23), which means that very close to theIF this is only about 10 years and most of the main emission

has a timescale of less than a century. Therefore, variations inthe UV luminosity or distance of v1 Ori C can produce changesin the nebula. If one simply accepts the Yale Bright StarCatalogue (Hoffleit 1964) value of 133 km s21, then v1 Ori Cis moving into OMC-1 at a rate of 6 km s21 and would travelto the present position of the main IF in only 50,000 years. Itis very difficult to accept such a radial velocity because it wouldhave a corresponding proper motion much larger than the othercluster stars (Jones & Walker 1988), and it is hard to understandwhy the most massive star in this cluster (Hillenbrand & Hart-mann 1998) could have such a large spatial velocity. This pre-cipitated a literature search by the author for various deter-minations of the radial velocity and an attempt to assess theiraccuracies.

Quite different radial velocities have been obtained by dif-ferent observers. Moore’s (1932) catalog of radial velocitiesgives km s21, which is similar to the less accurate23.1 5 1.8value of km s21 of Plaskett (1924). The latter gives24 5 16a value of km s21 for the Ca ii lines, whose average16 5 5modern value is about km s21 with the uncertainty20 5 4arising from how the complex profile of the absorption bymultiple narrow lines (O’Dell et al. 1993) would have beentreated in the measurement by eye of low-resolution spectra.Hoffleit’s value seems primarily driven by the study of Struve& Titus (1944), who measured all of the stars of the Trapezium.They found km s21 for v1 Ori C, km s2137.5 5 2.0 0.5 5 2.3for the nebula, and km s21 for Ca ii; i.e., their20.3 5 2.9Ca ii values agree with modern values, but the nebular valueis low by about 18. The latter difference raises fundamentalquestions about the small probable errors they claim. A morerecent high-resolution study by Conti (1972) found km26 5 3s21, while a lower resolution study by Abt et al. (1991) found

km s21. A program of monitoring the spectrum of17 5 11v1 Ori C at high resolution by Stahl et al. (1993) gave a valueof km s21, although a brief period of a significantly14 5 1lower velocity was seen. The author and his assistant DavidBrown measured the stellar absorption lines in the spectra ob-tained for studying the interstellar absorption lines (O’Dell etal. 1993) and found km s21. The most recent set of14 5 1observations by Stahl (1997) gives radial velocities with anaccuracy of a few km s21 over an interval of about 1900 days.During this interval he saw nonperiodic short-timescale vari-ations of about 10 km s21 and values through the range of

km s21. This indicates that the earlier observations were5–43not anomalous but are part of a currently not understood patternof change. Certainly there is not a monotonic secular change.

The explanation of the remarkable radial velocity behaviorof v1 Ori C probably lies with the intrinsic nature of the star.Conti (1972) established that the He ii l4686 line was asym-metric, being of a P Cygni type. Walborn (1981) first estab-lished that there were systematic variations in the spectrum ofv1 Ori C, a conclusion verified and extended by the first studyof Stahl et al. (1993) and additional observations (Stahl et al.1996). The object even seems to vary in its X-ray luminosity




2001 PASP, 113:29–40

(Gagne et al. 1997). All of the authors agree that there is a15.4 day periodic variation, although there is no short-termradial velocity variation. An oblique magnetic rotator model isusually invoked to explain the spectroscopic and X-raychanges, which would be consistent with demonstrating nodetectable radial velocity changes. Stahl (1997) favors an in-terpretation of the radial velocity changes as patterns of changein the profiles of the stellar absorption lines, although the lackof periodicity of the velocity changes can be used as a coun-terargument to this interpretation, leaving open the chance thatthe variations are due to the motion of the star. At this pointin time the important issue of the velocity of v1 Ori C remainsquite unresolved.

6. SCATTERED LIGHT AS A USEFUL ANDCONFUSING FACTOR

Scattered starlight dominates the spectrum of the continuum,being about 5 times stronger than the atomic component of thegas. This fact has been known from the early images of Wurm& Rosino (1956, 1957) and was one of the contributing factorsin Wurm’s argument for a slab model of the nebula. Thestrength of the continuum was also measured by photoelectricfilter photometry (O’Dell & Hubbard 1965) and then morerecently by Baldwin et al. (1991) in their slit spectra. An ex-tensive set of calculations by Schiffer & Mathis (1974) triedto explain the observations, but these models assumed that thedust was distributed close to the source stars, and are probablynot relevant.

The fact that starlight is scattered by dust in the nebula shouldnot be a surprise. We know that the visual optical depth behindthe main IF is very large, so that as long as the albedo of thedust particles is nonzero, there should be a measurable scatteredlight continuum. In effect, the facing side of OMC-1 acts as adiffuse mirror, a fact that was appreciated by Munch (1985)but published only in a publication that was not widely cir-culated. Many of the elements of his discussion of the effectof scattering of the He i l10830 line were independently in-troduced by the author (O’Dell, Walter, & Dufour 1992).

In his study of the [O iii] l5007 line, Castaneda (1988) madea multiple Gaussian component fit to each of his spectra. Inthis analysis he identified three strong velocity systems. SystemA was the strongest and represents emission near the IF. Hisdeconvolutions also yielded a second system shifted by just afew km s21 blueward, which was never fully resolved. Thistype of deconvolution into multiple components was made be-cause we were seeing for the first time multiple higher velocitysystems, and it was thought that the asymmetries in the mainemission were due to unresolved components. Henney (1998)has shown that this second component is artificial, simply rep-resenting the acceleration of gas away from the IF. However,the third systematic component cannot be explained in this way.This component is characteristically twice the FWHM of themain component and has a ratio of brightness of about 25%.

Its peak is also typically redshifted with respect to the maincomponent. The interpretation applied by O’Dell et al. (1992)is that this is the reflection of the main component, with theline being slightly redshifted because the light has been re-flected from a mirror that is moving away from the emittinglayer, which produces a doubling of the velocity shift in thereflected light (of course it is actually the emitting layer thatis moving away from the stationary OMC-1). They interpretthe broadening as being caused by light being reflected (scat-tered) from a variety of relative velocities, since both the scat-tering and emitting layers are extended. Light scattered fromdirectly below the emitting point will have the greatest redshift.Detailed models of this process are included in a recent paperby Henney (1998).

This is not simply a second-order effect. Independent of ourunderstanding of the details of the blister model, this modelunavoidably says that a certain significant fraction of the ra-diation emitted in the direction of OMC-1 will be backscatteredtoward the observer. This factor must be included in our in-terpretation of all facets of the nebula. For example, the cal-culation of an equivalent thickness of the emitting layer fromthe known surface brightness and electron density assumes thatthe emission occurs isotropically into 4p steradians. In the caseof the dust being a perfect reflector (albedo unity), then theemission is radiated into 2p steradians. Without knowing thedetailed processes involved with this scattered light, it is notpossible to make accurate corrections. The problem is evenmore complex, since the fraction of light that is backscatteredprobably depends upon wavelength and the separation betweenthe emitting and scattering layers. The wavelength dependence(which would arise from the effective change in optical depth,albedo, and phase function with wavelength) would mean thaton the average a different fraction of the light is scattered atdifferent wavelengths. If this fraction increased with decreasingwavelength, then it means that the physical interpretation ofthe observed spectrum would be skewed by artificially en-hancing the bluer lines. If the blue lines are artificially en-hanced, this could provide an explanation of the the anomalousextinction law for M42 found by Costero & Peimbert (1970),although this would not remove the peculiarity of the Orionextinction law found by investigations of the stars lying in frontof the nebula. The case becomes even more complex if thefraction of light that is scattered depends upon the separationof the emitting and scattering layers. This separation would bequite small for ions such as [O i] and [S ii], which arise fromclose to the IF, while it would be large for high-ionizationemission from [O iii]. This correction introduces a new levelof complexity into analysis of spectrophotometry of M42 andmust be causing a misinterpretation of the observed lines. With-out understanding the details of this process, we must assumethat errors in the relative values of emission lines can be asmuch as 25%.



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7. MAJOR UNRESOLVED QUESTIONS ABOUTTHE STRUCTURE OF M42

Although there have been many advances in our knowledgeof the structure of M42 since the Henry Draper Symposium,a number of fundamental issues remain unresolved. This clos-ing compilation is basically a list of the ones that are the au-thor’s favorites, and their solution should help clarify what isactually happening in this important object.

7.1. Why Is the Bright Bar So Linear?

It is easy to cause small-scale linear features, with jets andshadows being obvious mechanisms; but these cannot be in-voked to explain this major feature in M42. Clearly the BrightBar is a warp in the main IF of M42 which we view fromnearly edge-on, but this means that the radiation field that hassculpted this part of the nebula has been constant over a longlength or that the underlying density of neutral material isnearly constant.

7.2. Why Can We See Such Fine-Scale Structure?

The main piece of physics that describes the brightness ofthe material near the main IF is the fact that the local surfacebrightness will depend only on the flux of ionizing photonsfrom v1 Ori C and the viewing angle. The surface brightnesstogether with knowledge of the emitting region densities givethe result that the e21 scale for emission is 140. This meansthat we would expect the nebula to be smooth at this scale,whereas the observed nebula shows structure down to a fewarcseconds. The explanation cannot lie with density differencesof the material, because if the object is ionization bounded, thesurface brightness depends only on the ionizing flux. Perhapsthe answer has its source with density in the sense that theadvance of the IF into OMC-1 will proceed more rapidly wherethe density is low, leaving behind columns of un-ionized ma-terial (the much publicized pillars seen in M16 are the mostdramatic example, which has been known since at least thetime of the Palomar Sky Survey). Irregularities in OMC-1 atthe scale of 0.01 pc would then show up as roughness in M42’smain IF, which would become visible through variations of thelocal flux and viewing angle. There are several arguments thatsuch irregularities exist in the PDR (Hollenbach & Tielens1997), as high-density knots are necessary to produce the COmillimeter line profiles (Stacey et al. 1993).

7.3. What Is the Source of the Extra Line Broadening?

It is remarkable that all of the emission lines studied to datehave an unexplained line-width component of about 9 km s21,this being true for the [O i] emission that must originate exactlyin the IF and the [O iii] that must arise from a more extendedzone closer to v1 Ori C and farther from the IF. It is as ifionizing gas is moving in random directions at about this ve-locity on a scale of less than a few arcseconds. What physical

model can explain this and possibly what nebular physics isbeing left out of what we are doing in building these models?If the timescale for driving this extra line broadening is short,then this is an important element in the determination of theenergy balance of M42. The unexplained line broadening maybe produced by the same small-scale structure discussed in theprevious section. In this case the broadening would be causedby free flow of material in all directions from around isolatedknots and pillars in the IF (O’Dell & Wen 1992). An attractivealternative is that of Ferland (2001), who invokes the presenceof Alfven waves.

7.4. How Is Scattered Emission-Line Radiation AffectingAnalysis of the Spectra?

We now know that not only is the visual starlight from theTrapezium stars scattered by the concentration of particles im-mediately behind the IF in the associated PDR, but that thisscattering also occurs for the emission lines. Short of observingevery line of interest at velocity resolutions sufficient to de-convolve the original and scattered components, this mecha-nism introduces a nonphotometric “noise” into the analysis ofemission-line ratios. Many of the problems dealing with in-terpretation of the spectrum of M42 will not be changed byphotometric uncertainties at the level of 25%, but others will,and this needs to be considered in all spectral analysis.

7.5. Is There a Central Cavity?

We certainly know that v1 Ori C is the source of a strongstellar wind (Howarth & Prinja 1989), which raises the ex-pectation that there would be a central cavity of highly shockedmaterial. This has been invoked by many of the early modelsof M42 and as recently as the paper by Wilson et al. (1997)to explain the more redshifted velocities near the center of thenebula. However, the author favors an interpretation of thatfeature as being due to radiation pressure constraining flowaway from the main IF in that region. Furthermore, there is noevidence for the high-temperature shocked material that wouldbe formed behind the shock that leads this wind. The mostlikely explanation is that the flow of material through the IFand into the region that produces the visible M42 is energet-ically much more important than the stellar wind of v1 Ori C.This is supported by the fact that the energy flux in particlesis ergs s21 while the energy flux of ionizing photons351.3 # 10is ergs s21. There certainly are large-scale structures381.7 # 10in the central part of M42, but none of these are centered onv1 Ori C (O’Dell et al. 1997a).

7.6. Where Are the Non-Lid Absorption Lines BeingFormed?

The riddle of the location of the material forming the He iabsorption lines has been with us for half a century withoutresolution. We have now added the problem of also explainingNa i and Ca ii lines at both blueshifts and redshifts. There are




2001 PASP, 113:29–40

sources for pushing nebular material around at these velocitiesin the form of the outflows from pre–main-sequence low-massstars, but at this point we do not have the theoretical knowledgethat links the two.

7.7. Where Is the Lid Located?

The presence of a lid seems well established, being seen inH i, Na i, and Ca ii absorption, indicating that there must bea second IF on the side of the Lid facing v1 Ori C. The importantquestion is one of how far from v1 Ori C it is located. Theanswer has broad application not only in study of the nebulaproper but also in study of the Trapezium cluster, as a close-in Lid would indicate that some of the observed cluster starsare probably located within or on the near side, while most liewithin the photoionized main cavity.

It is the author’s pleasure to acknowledge the independentwork, assistance, and discussions of his former students Hector

O. Castaneda, Michael R. Jones, Leisa K. Townsley, ZhengWen, Xihai Hu, and Wendy M. Lane in formulating the viewof M42 that appears here. My more senior colleagues GaryFerland, Patrick Hartigan, Will Henney, and John Meaburnhave also been generous of their material and ideas over thelast many years. This review is based on an oral paper presentedat the 1997 June Ringberg Castle conference on “The OrionComplex Revisited” sponsored by the the Max Planck Society.The original draft of this review was done while supported bythe German Alexander von Humboldt Foundation and hostedby the Max Planck Institute for Astronomy in Heidelberg, Ger-many. Final preparation of the manuscript was supported inpart by the Space Telescope Science Institute grant GO-8121.

The figure depicting a close-up view of M42 was directedby Carter Emmert, with technical production by Erik Wesselak,both from the American Museum of Natural History in NewYork. Volumetric rendering software was done by Dave Nadeauand Jon Genetti of the San Diego Supercomputer Center. Theirprofessionalism and cooperation are gratefully acknowledged.

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