molecular shocks in star forming regions

M O L E C U L A R S H O C K S IN S T A R F O R M I N G R E G I O N S

P. W. J. L. BRAND Institute for Astronomy, University o] Edinburgh, Royal Observatory Edinburgh EH9 3HJ

Abstract. The role of excited molecular hydrogen as a powerful observational tool is examined in the context of shock phenomena in molecular clouds, particularly in star forming regions. Conclusions that may be drawn from line intensities and line profiles, and the properties of J and C shocks in a bow shock structure are discussed.

Key words: Molecular hydrogen - shocks - outflows

1. I n t r o d u c t i o n

The realization that forming stars produce powerful supersonic molecular wind, often collimated into bipolar jets, came from mapping CO line profiles (cf. Lada 1985).

Soon afterwards the discovery of shocked molecular hydrogen in Orion (Gautier et al. 1976) led a series of papers describing shock models. The observation (Nadeau & Geballe 1979) of wide line profiles confirmed the dynamic origin of the emission, but raised problems for the models which have not yet been satisfactorily overcome. This investigation of shocks using molecular hydrogen lines which is still in progress~ is the subject of this arti- cle. The remainder of this paper will develop current understartding of shocks by considering evidence from Orion (section 2), Herbig-Haro 7 (section 3) and other sources (section 4). Section 5 lists briefly the current situation and expected progress.

The theory of interstellar shocks has been recently reviewed by Dralne & McKee (1993) and that review contains important insights into the physics of molecular shocks, and references thereto.

2. O r i o n

2.1 PREAMBLE

The shocked outflow region in the Orion molecular cloud just behind the Trapeziumdit HII region was mapped by Beckwith et al. in 1978.

The first line ratio measurements (Beckwith et al. 1979) showed that the line emission was thermalized, and hydrodynamic shock models (HoHenbach & Shull 1977, Kwan 1977, London et al. 1977) were able to explain the H2 emission. It was quickly realized (Kwan ibid.) that since H2 was completely thermally dissociated by a non-magnetic shock travelling faster than 22 km s -1 through the gas, the wide velocity profiles observed by Nadeau & Geballe (1979) were a problem for the models. The introduction by Dralne (1980,

Astrophysics and Space Science 224: 125-138, 1995. © 1995 Kluwer Academic Publishers. Printed in Belgium.

126 P . W . J . L . BRAND

also Draine, Roberge & Dalgarno 1983) of 'C shocks', which are able to travel at >40 km s -1 through H2 without completely dissociating it, led to several models that were consistent with the contemporary measurements of H2 line intensities, and also gave a better prediction of the amount of hot CO (Draine & Roberge 1982).

By 1988 (Brand et al. ) the H2 line ratio measurements for Orion had been extended sufficiently to rule out the plane parallel C shock models. Relatively too much emission was observed from the lines of highest excitation to match the models. A brief discussion of model properties, needed to demonstrate the differences, follows.

2.2 J AND C MODELS

Draine (1980) used the term J shock (jump shock) to denote the hydrodynamic jump followed by a cooling zone which constitutes a shock in a non-magnetic medium. In contrast, the C shock (continuous shock) arises in a magnetised medium consisting predominantly of neutral material mixed with a small population of ions. These ions are taken with the magnetic field to be a magnetohydrodynamic fluid which supports Alfv4n and magne- tosonic waves. The neutral fluid interacts with the magnetised fluid through ion-neutral collisions, which produce dynamic friction and heating of the neutral fluid. A plane pressure wave travelling through the magnetised fluid perpendicular to the field lines drives a submagnetosonic compression ahead of it into the undisturbed magnetised fluid, and the resulting flow of ions through the neutrals heats them and accelerates them in the direc- tion of propagation. The disturbance is submagnetosonic with respect to the ion fluid because of the low ion density, but is supersonic as far as the neutral fluid is concerned. If the coupling between ions and neutrals is suffÉ- cient, the neutral fluid will have been heated and accelerated sufficiently by ion/neutral drag to achieve its 'shocked' condition (required by conservation of momentum and mass) without passing through a hydrodynamic 'jump'. This structure is called a C shock, extensively investigated by Draine et al. (1983).

The properties of these two shock types are contrasted in Fig. 1. If the post-shock density is high enough, then the gas will be locally in thermody- namic equilibrium. This appears to be the case in most of the shocks which have been observed.

The temperature profile of the neutral gas is of particular interest. In the J shock it rises virtually immediately to its maximum value, and then drops as the gas cools, under nearly isobaric conditions. The rate of cooling is determined by the cooling function of the neutral gas.

Then the column density of each level of excited H2 in the cooling zone will depend on the run of temperature with column density in the shock model and the ratio of these cohmn densities will depend only on the cooling

MOLECULAR SHOCKS IN STAR FORMING REGIONS 127

&

.%

b

Fig. 1. T profile of (a) a J shock and (b) a C shock. The dashed line denotes velocity of the neutral gas, and the dotted line denotes ion velocity.

function. Cooling at the highest temperatures in these shocks in molecular clouds is due to H~ dissociation, and causes complete dissociation in shocks travelling faster than 22 km s -1 through dense gas (Kwan 1977, Brand 1992). At lower temperatures (< 3500K) H2 line cooling and then cooling by molecules dominates.

In contrast, the run of temperature of the neutrals in a C shock is a smooth function, rising to a maximum close to the region where drag is greatest, and then subsiding in a cooling zone. The temperature profile can be roughly approximated by a ' top-hat ' distribution, whose peak value is determined by the frictional heating rate (which depends, in ter alia, on the velocity). Thus the H2 emission in this case will look more like that from a slab of gas at a uniform temperature set by the shock parameters, but not depending on the cooling function.

Fig. 2 shows the run of column density in H2 versus energy of the level predicted by a naive J shock. The concave curve is characteristic, mapping the cooling rate as a function of T in such a shock. A constant T distribution would be represented by a straight line in such a plot (for example the 2000K line indicated in the figure by a dashed line).

Relative to such a straight line, the J shock model predicts excess high and low excitation emission. This is emphasised in Fig. 3 which displays the results from Peak 1 in the Orion outflow. Here the plot has been skewed by dividing column density values by those from a Boltzmann distribution at 2000K to accentuate those departures from a Boltzmann distribution. The data are well-fit by a naive J shock model. The dashed lines represent the predictions from plane C shock models (Draine & Roberge 1982, Chernoff et al. 1982), and these, having the expected flat distribution, are ruled out

128 P . W . J . L . BRAND

0 v

:Z

z-

i 0 z°

i 0 "

10 '8

I 0 "

i 0 ' 8 i

I0 's

i0 '4

i0 '3

I 0 '~ 0

E 6 -~ E ock n=10 cm , V,=I0 k m / s -=

, , , i I , , , , I , J , , l , , , , I , , , ,

5000 10000 15000 20000 25000 Energy Level (K)

Fig. 2. Post-shock cooling: logarithm of column density in each energy level of He versus energy (after Burton (1992).

lOO

lo

1

lO 20 3O

Energy of upper level ( 1000 K )

Fig. 3. Logarithm of column density in each energy level versus energy, inferred from observations of Peak i in the Orion molecular outflow.

by the data. Alas, that is not the end of the story•

2.3 VELOCITY PROFILES AND BOW SHOCK STRUCTURES

The velocity profile of a bright H2 line at the same position in Orion was measured to have a line width of 140 km s -1 FWZI (Brand et al. 1989a) which cannot be produced in a single plane J shock, with its constraints on velocity. In the first place, the shock has a maximum velocity of 22 km s -1 before complete dissociation. In the second place the velocities predicted by J shocks for H2 lines are all close to the final gas velocity, producing a very small range of velocities. In the case of a C shock the limits are relaxed,


~V(shock)

Fig. 4. Bow shock structure used to explain the H2 spectra.

but neither the increased maximum (to 40 km s -1) nor the possible range of velocities will be sufficient to explain a profile such ~ that at Peak 1 in Orion.

The observation requires for its explanation a large range of shock speeds in the same small region of sky which is observed. The profile was obtained in a 5 arcsecond aperture (with the UK Infrared Telescope and its CGS II spectrometer) - - a projected size of 3x1016 cln at Orion. The easiest way to meet this tight constraint is by modelling a bow shock in the molecular material (Fig. 4).

In this structure, gas flows through the curved surface, with a variety of normal velocities and therefore of shock strengths. Provided the structure is small enough, the observed emission will be the sum (or rather the integral) of all of these differently directed shocks. How does this change the plane shock predictions? This has been well-treated for atomic emission by Harti- gun et al. (1987). They demonstrate that a range of velocities equal to the velocity of the structure through the gas is seen in line profiles taken from the whole shocked region and demonstrate how to infer this velocity and the angle with respect to the line of sight of the supersonic flow.

Fig. 4 shows the model used to investigate the line profile. The shape of the bow shock is prescribed (in fact it is determined by the behuviour of the shocked gas and its interaction with whatever is driving the shock, but this is considered still to be ill-determined). The shock structure may be considered to be driven by a jet or a bullet through stationary gas. This is not critical for the model - - a simple velocity shift can be applied to allow for a 'wind hitting an obstruction' scenario. The shocked gas is given an impulse perpendicular to the shock surface, and radiates line emission into the corresponding part of the line of sight velocity profile. The velocity

130 P . W . J . L . BRAND

profile of a line is then the integral of each of the elements of shocked gas over the shock surface, each element with its own line of sight velocity.

Molecular gas near the head of the structure will meet the shock more or less head on, and (assuming a high velocity for the gas through the shock structure) the molecules will be dissociated by the shock. Only further down the flank, with a more oblique shock, will molecular gas survive and radiate.

The importance of this model is obvious. It provides the possibility of explaining the wide velocity profiles. It also re-opens the issue of line intensity ratios from J and C shocks. It will be recalled that plane J shocks naturally reproduce the curvature seen in the plot (Fig. 2 or Fig. 3) of log (column density) versus energy level for H2 line emission, whereas plane C shocks tend to produce straight (Boltzmann-like) plots. However, a moment's con- sideration will show that the sum of two or more differently excited plane C shocks will produce a concave plot! But by hypothesis we are looking simul- taneously at shocks of different strengths when we observe a bow structure. Thus, we cannot now choose between J and C shocks on the basis of this property.

2.4 EXCITATION CONDITIONS

There is one further result from the Orion shocked region which bears on the problem (Brand et aI. 1989b). This is shown in Fig. 5, and is a plot of the ratio of two lines (0-0 S(13) and 1-0 0(7)) which have nearly the same wavelength (and so are not subject to differential extinction), but originate from different energy levels. Thus, their ratio determines the particular curve of column density versus energy level at the point of observation. The plot is of this ratio versus line intensity, from points all over the region originally mapped by Beckwith et al. (1978).

It is clear from Fig. 5 that the ratio does not vary significantly, and that therefore it is independent of position throughout the map. Apparently all of the shocks in this region have the same excitation state. What does this signify?

In the case of plane J shocks, it implies that the cooling curve is the same everywhere, but is insensitive to other shock properties. Only if all cooling processes have the same density dependence will this condition be met. This would be so if the important coolants have the same density dependence. But for instance H2 dissociational cooling per unit volume depends on density squared, while H2 line radiation cooling depends (at these densities) on density linearly. Thus, if as expected these are major contributors to the cooling function, different strength shocks will have different cooling curves, and it becomes hard to claim that plane J shocks can produce the result in Fig. 5. (However, a recent result by Smith (1994) which is disctlssed in section 4.4 may modify this conclusion).

A fortiori, C shock excitation is strongly dependent on velocity and so


1.5

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*," .÷ ~ . . . . , ~*÷,÷, *;: ¢ 4 + + "44 + " 4 + 4 4

* ++ 4" +÷

0 ' ' ' ' J . . . . f . . . . 0 10 20 30

0 ( 7 ) + S ( 1 3 ) / ( t O - 1 8 W m - 2 b e a m - ! )

Fig. 5. The ratio of the H2 lines 0-0 S(13) to 1-0 O(7) as a function of line intensity in the region of the Orion outflow mapped by Beckwith et al. (1978).

the constancy in Fig. 5 implies all shocks have the same speed at least - - an even t ighter constraint.

Now, as pointed out by Smith et al. (1991a) the bow shock structure provides a potential cure. The leading part of the bow structure consists of a completely dissociating shock, and only further down the structure does It2 survive to radiate. The range of excitation observed is that arising from between this survival zone and the back of the structure where the shock is too weak to excite the observed H2 lines. This region, which I will call the excitation region, exists where perpendicular shock velocities range from Vmax to 10 km s -1 (the velocity necessary to strongly excite the 1-0 S(1) line). The velocity Vm~x for C shocks may be 50 km s -1 (or greater: cf. Smith et al. 1991b). For any given incident velocity of gas on shock structure - - bow shock speed - - typically up to several hundred km s -1 , the excitation region will encompass a range of incident perpendicular velocities. This range will remain the same for different bow shock speeds although the excitation region will move further aft as the bow shock speed is increased. The contribution to intensity and velocity profile will be determined by the bow shock shape in the excitation region.

Then different velocity structures will give the same excitation, provided that the shape of the bow structures are all the same and approximately paraboloidal, in a nearly scale-free manner. This seems an easier constraint but (Smith et aI. 1991a) it requires that we observe the entire bow structure, and that certain other conditions are met.

132 P . w . J . L . BRAND

2 .5 HIGH MAGNETIC FIELD BOW SHOCKS, BULLETS~ AND OTHER MAT-

TERS

The ability of bow shock structures to explain not just the excitation but also the velocity profiles of H2 which are observed depends on the ability of such a structure to turn the emitting gas in a variety of directions.

It was shown (Brand et al. 1989a) that non-magnetic J shocks are unable to do so effectively enough. Even without the tight constraint on maximum velocity before complete dissociation (22 km s- l ) , the profiles carry too little power in the wings to replicate the observations. The problem lies in the ratio of observed velocity width (and hence total structure velocity) to the velocity corresponding to maximum emission of H2. If this is high, then that fraction of the gas in the flow that contributes significantly is turned through a small angle, and will therefore not contribute much to the profile outside velocities of one or two tens of km s -1.

C shocks can turn the gas without destroying the H2 because the Alfvgnic Mach number is relatively small if the magnetic field is high. Smith et al. (1991b) showed that both the line ratios and velocity profiles of H2 can be replicated in Orion (Peak 1) by a bow-shock structure in which the shock transition is by C shock, and the magnetic field is particularly - - but perhaps not unrealistically - - high.

There are several lingering doubts. In the first case, the C shock is oblique, and as shown by Wardle & Draine (1987) and Wardle (1991) such shocks can be unstable.

In the second case the constraints mentioned earlier still apply, namely the entire flow is required to contribute to the line intensities and profiles, a~d the shape and cooling function requirements are rather tight. Thus the situation is still far from clear. Indeed Smith (1994) has shown that a highly magnetized J shock is still a possible contender.

A remaining possibility is that proposed by Graham et aI. (1991,1992) for H2 emission seen in the Cygnus Loop, and by Hartigan et al. (1989) and by Cart (1993) for emission from Herbig-Haro 7. These authors suggest that the emission is due to a magnetic precursor (in effect, the leading edge of a C shock) followed by a standard J shock which completely dissociates the molecular component. This possibility has not been fully explored, aad has several attractive features, but doubt as to its stability against rupture to a J shock and its ability to produce sufficient column density have been voiced (e.g. DraJne & McKee 1993).

A recent and exciting development has been the detailed confirmation by Allen ~z Burton (1993) that the faint H2 'fingers' of emission seen to the north-west of the Orion outflow (Taylor et al. 1984) are bow shocks driven probably by bullets originating at or near the source IRc2. The iron emission seen at the tips of these bow shock structures is consistent with a model in


which the H2 wakes are the tails of the shock structures. Work is in progress (Chrysostomou et a l . , Tedds et aI. ) to decide if a

model based on an ensemble of such objects will explain the entire flow, or if these 'bullets' are the exception in a wind or blast that inflates a 'bubble' of gas inside the molecular cloud.

3. Herb ig -Haro 7

We turn now to Herbig-Haro 7. This object is clearly the bow shock at the end of a jet apparently originating in the young stellar object SSV 13. Maps by Hartigan et al. (1989) show the H2 emission surrounds the ionic emission associated with the jet, and the investigation by Cart (1993) shows the velocity structure in the H2 profile map with the characteristics of a bow shock.

The bow structure is resolved (5 arc seconds) and can be used to test predictions of the bow shock models proposed to explain the Orion results.

If any dust in the structure is assumed to be optically thin at the H2 lines (this is consistent with Carr's measurement ibid.), then an arbitrary orientation of the bow shock will imply that each line of sight intersects the structure in two points, with differing shock strengths. Thus as noted at the end of subsection 2.3, a curved distribution of log (column density) versus energy level will be seen.

The outcome of the measurements (Fernandes et al. 1995) exceeds expec- tations. The resulting column density plot (Fig. 6) shows the data for column densities, and a good fit of a model consisting of a bow C shock plus flu- orescence, the latter assumed to be due to the high excitation of the bow at its head. If this is fully borne out, then the simultaneous conditions this imposes on the model will lead to a rather exact description of such bow shocks, and the implied UV field - presumably from the hot head of HH7 - will set a lower limit on the degree of ionization in the vicinity which will have implications for magnetic precursor models (and indeed for C shocks themselves).

There remains the possibility that what is being observed is not fluores- cence but rather the re-formation of H2 after dissociation on passage through the shock. A decision on this awaits a fuller data set, and more precise modelling. In any event, identification of the source of the non-thermalised lines remains a top priority.

4. E l sewhere

The data on column densities and line profiles in Orion and Herbig Haro 7 have reached an exciting juncture. But what of other evidence bearing on the J shock/C shock debate?

134 P . W . J . L . BRAND

I 00

. . . . I . . . . I . . . . I . . . . I . . . . | • , ,

X J~ ,wC .dx=k+F l , . , ~ ,~ , * ( ~ ' ~=~ r . s . [ ~ }

• ~ .0~ 7 ( ' 3 " ) mw 34

. . . . f . , , , f , , . . f . . . . I • . . , I . . . .

Uppc¢ k~ ! ~ y ( I { ¢ I . t Ln )

Fig. 6. Normalized line H2 column density ratio versus energy level in HH7 (from Fer- nandes 1995)

I discuss briefly the evidence from a massive outflow DR21, two supernova remnants IC443 and the Cygnus Loop, and two jet-like outflows from young stellar objects HH90A and L1551.

4.1 DR21

This is possibly the most massive outflow source. The hydrogen emission delineates a collimated bipolar flow from a heavily obscured HII region, the emission being highly d u m p e d (Garden et al. 1990) in both lobes, the clumps having a range of velocities and velocity widths. In the east lobe the lines are relatively narrow (<20 km s- l ) , while in the west lobe velocity widths of up to 100 km s -1 are seen (Garden et al. 1991). Measurements of line ratios in this region (Fernandes et al. 1995) have been made and show, over the limited area that has been investigated, that the emission is explained by C shock excitation, but with a small component of fluorescent emission, as in HHT.

4.2 IC443

This extended supernova remnant is associated with a molecular cloud in which the expanding blast wave can be seen (Burton et al. 1988) as a ring of shocked H2. This shock is also seen in HI (Braun & Strom 1986) and CO (DeNoyer 1979, White et al. 1987).

Measurement of H2 column densities (Moorhouse et al. 1991) show the characteristic shock signature, and when originally interpreted as a/l shock, indicated that multiple shocks had to be involved. Burton et al. (1990) find that the [OI] 63#m emission from this region is consistent with C shocks provided special constraints apply to their velocities, or with J shocks pro-


vided that H20 and OH are not present. Neither explanation is satisfying. Interferometric CO observations by Wang & Scoville (1992) of the region lead them to model the CO dynamics by a doubly shocked cloud as did Moorhouse et al. (ibid.) but with a one shock of C type and the other, the incident supernova blast wave, a J shock. The situation for the IC443 molecular shock is still unclear.

4 .3 CYGNUS L o o p

A very surprising result by Graham et al. (1991) demonstrates that in the north eastern filamentary optical nebulosity there is also shocked HI2 emission. This appears on the outside of the Ha filaments, and so cannot be the strongly cooled post-shock gas. The explanation given is that a magnetic precursor is being observed, and indeed it is hard to envisage any other simple explanation. But as already noted at the end of Section 2 there are concerns about the stability of such structures, especially in a region where ionization is not necessarily small (cf. also Elmegreen & Combes 1992).

4.4 HH91A

Gredel et al. (1992) measured many H2 lines in this complex part of an outflow which excites HH90 and HH91. HH90A is a bright knot very close to an infrared source whence the flow possibly originates. The remarkable fact about the H2 line intensities is that the derived plot of log (column density) versus energy level is straight, corresponding to thermal excitation at T=2750K. From what has been said, it is clear that this cannot arise from J shock excitation, and if it arises from a C shock, it must be a single plane shock. There is a rapid change of intensity with position in I-IH91A and the dominant component is very small. It is hard to envisage a geometry (possibly spherically symmetric) which would produce such a remarkable column density distribution.

It may be that since the highest energy level observed was at only 14500K, curvature is present and indeed the lowest levels deviate very slightly above the line. A C shocked bow viewed with special geometry might then be a possible explanation. However, new J shock models by Smith (1994) can fit these data. If this is confirmed, this most intense source will be a severe test for theory.

4.5 L1551

As a last example of data which may enable a clearer view of molecular cloud shocks, the CO interferometer map of L1551 (Barsony et al. 1994) shows several filaments which are asserted to be shocks, and which coincide with shock-excited Ha and [NII] emission (Graham & Heyer 1990). In the brightest of these, a smooth gradient of velocity of 8 km s -1 over 1016 cm is observed and might be the signature of a C shock.

136 P . W . J . L . BRAND

The range of velocities (8 km s -1) is too s m a l l - unless projection effects are significant - - to give strong H2 emission, but this source may be an important testbed given its apparent planarity over 2 arcminutes.

5. T h e ups a n d downs of shocked I-I2 m e a s u r e m e n t s

A review of where we are and where we are going in the attempt to under- stand shocks in star-forming regions turns out to be a mixed bag, and I list here progress points: past, present and future.

a) The jury is still out on J versus C shocks as the exciting mechanism for most H2 emission. The weight of evidence is swinging towards C shocks, but (e.g. Smith 1994) J shocks will not go away. The best data sets (Orion, HH7 and HH91A) each appear to tell different stories.

b) It seems clear that bow shock structures are necessary to produce the nearly omni-present broad velocity profiles. But this explanation requires very high magnetic fields (Smith et al. 1991b), and observing apertures big enough to include the entire structure. For the nearest sources, such as Orion, this latter requirement seems to be a tall order, and the constancy of the ratio of column densities in Orion (Brand et al. 1989b) tightens further this constraint.

c) Perhaps the most critical question about the structure of these shocks is that of cooling functions. In particular, is water the dominant coolant below 1000K? This question has been hard to answer observationally to date. It is one of the more important goals of the ISO mission to find the a n s w e r .

d) The first steps are being taken to produce detailed models of detailed observations of shock structures. For the nearest sources maps with a resolution of 1 ~ c second or better, and velocity resolution of the order of 10 km s -1 should constrain the models so that the J or C pips squeak. The fortunate fact is that this capability is now becoming available. It is one of the pleasures of this research that as new questions arise, so do the means to seek the answers.

e) In similar vein, the cooling lengths in the shocks being discussed range f r o m 1013 cm for J shocks up to 1017 cm for some C shocks. This translates to a few milli-arc seconds to several arc seconds in the nearest sources. Exam- ination of the caustic edges of shock structure images with the enhanced resolution becoming available with adaptive optics should give a clear dis- tinction between J and C shocks.

Thus, on the one hand it is frustrating how slow progress has been in the development of our understanding of shock physics in star forming regions. We still do not have a clear idea of how these things work.

On the other hand, current and future observational capabilities are exactly what are required to address the now much clearer questions about


the shock mechanisms. I am moderately confident that in three or four years time this subject will have given shock theory the diagnostic power needed to measure the conditions in regions where stars are disrupting their birthplaces.

Acknowledgements

Thanks are due to colleagues who helped to form these ideas, particularly Michael Burton, Tom Geballe, Michael D. Smith, and also to the staff of the United Kingdom Infrared Telescope, a facility of the Particle Physics and Astronomy Research Council.

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