1. Hydroxyl Emission from Shock Waves in Interstellar Clouds Catherine Braiding

2. Hydroxyl Emission in Interstellar Clouds Supernova Remnants + Molecular Clouds OH Masers Shock Waves OH emission from Shock Waves Modelling OH Testing the Model Modelling Shock Waves Future Directions

3. Molecular Clouds About half the gas in the Galaxy is found in clouds of dense gas. These are cold enough (10-30 K) to form molecules. Gravitational collapse causes star formation. The clouds are dispersed by ultraviolet radiation, stellar wings and supernovae.

4. Supernova Remnants

5. Supernovae Mark the death of massive stars (>8M sun ). Distribute energy and heavy elements into the interstellar medium. Frequently occur near molecular clouds, due to the short lifespan of massive stars. Cause shock waves to be driven into the molecular cloud.

6. Supernovae + Molecular Clouds Wardle and Yusef-Zadeh (Science, volume 206, 2002)

7. Supernovae + Molecular Clouds Shock waves create compression and heating in the cloud. This can lead to star formation. The chemical composition of the gas is changed, as reactions between molecules are allowed to occur. It is difficult to positively identify this behaviour.

8. Supernovae + Molecular Clouds A “signpost” of the interaction is the OH 1720 MHz maser. About 10% of supernova remnants possess maser spots. By studying the emission and absorption of other OH lines in shocked gas as well as the maser spots, can gain a better understanding of the interaction.

9. OH Masers Microwave Amplification of Stimulated Emission Radiation: –Microwave analogue of a laser. –Occur naturally in stellar atmospheres and interstellar space. Bright, compact spectral line sources. These occur at 1612, 1665, 1667 and 1720 MHz

10. OH 1720 MHz Masers Not found in stellar atmospheres. Require specific physical conditions: –Density: n ~ 10 5 cm -3 –Temperature: T ~ 50–100 K –OH column density 10 16 – 10 17 cm -2 –The absence of a strong far-infrared continuum. Collisionally-pumped by H 2

11. (Pavlakis & Kylafis 1996, ApJ, 467, 300) OH Level Diagram

12. Shock Waves These conditions are satisfied if the shock is a slow, continuous shock wave. The low ionisation level in the molecular cloud causes the magnetic pressure to exceed the thermal pressure by several orders of magnitude. When a slow shock passes through, the ions stream ahead of the shock wave in what is known as a magnetic precursor.

13. Shock Waves *image of J vs C type shocks*

14. Shock Waves In C-type shocks, ion-neutral collisions smooth out the viscous transition, so that an extended region of gas is heated. Critical velocity for C-type shocks is 40-50 km s -1. Supernova-driven shock waves travel at ~25 km s -1.

15. Shock Waves All of the OH produced within the shock at temperatures above 400 K is converted rapidly to water. O + H 2  OH + H OH + H 2  H 2 O + H

16. Shock Waves The dissociation of water by ultraviolet radiation creates OH. H 2 O  OH + H X-rays from the supernova and cosmic rays induce a far-ultraviolet radiation field that is capable of dissociating water.

17. Shock Waves How does one identify these shocks? –OH 1720 MHz maser “signpost” –OH also detected in absorption –Known to be strong sources of H 2 2.12 µm emission –Contrast between CO emission in the both shocked and unshocked regions of the cloud

18. Candy (G349.7+02) – H 2 J. S. Lazendic et. al. in preparation

19. Candy (G349.7+02) – OH J. S. Lazendic et. al. in preparation

20. Modelling the OH Emission Wardle (1999) showed that by including photodissociation in the oxygen chemistry, the OH column density produced was sufficient to form OH 1720 MHz masers. This effect has not been examined in previous models.

21. Oxygen Chemistry in a C-type Shock (Wardle, ApJ, 525:L101, 1999)

22. Modelling the OH Emission Want to calculate the populations of the excited levels of OH for a given gas density, temperature and column density. Using this information, can then determine emission from one point in the gas. This can then be incorporated into shock calculations.

23. Calculating the Level Populations The level populations change over time as: *equation* These equations are integrated over a long period of time, so that many collisions and radiative transitions may occur, bringing the system to equilibrium.

24. Calculating the Level Populations Data was provided for the Einstein A coefficients for the first 32 hyperfine-split levels of OH. Given the high temperatures found in shocked gas, more levels were required for the model.

25. (Pavlakis & Kylafis 1996, ApJ, 467, 300) OH Level Diagram

26. Calculating the Level Populations The HITRAN 96 database contained level energies for the first 100 split levels of OH. Unfortunately, it only contained rotational transitions from the first 72 levels. However, the code can easily be updated when more data comes to hand.

27. Calculating the Level Populations The collisional rates used were obtained from Offer, Hemert and van Dishoeck, for transitions between the lowest 24 states. For the higher states, hard sphere rates were used.

28. Testing the Level Population Code For low temperatures and densities, the level populations should be concentrated in the lower levels. In the limits of high temperature or density, the population distribution tends towards a Boltzmann distribution.

29. Testing the Level Population Code *insert picture here*

32. Future Directions The shock code needs to be optimised for better runtimes. The calculated emission needs to be tested. The dependence of the emission on the input parameters will be explored. The effect of the X-ray flux on the emission should be examined.

33. Future Directions Calculations of the emission should then be compared with observations.

34. Future Directions Further observations of supernova remnant / molecular cloud interactions would provide greater opportunity to test this theory of OH emission. The GREAT spectrometer on SOFIA will be capable of detecting the warm OH column density within C-type shocks.

35. Future Directions SOFIA will fly in 2004 (we hope). (http://sofia.arc.nasa.gov)