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What Physical Processes Drive the Interstellar Medium in the Local Bubble?

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What Physical Processes Drive the Interstellar Medium in the Local Bubble? Space Sci Rev (2009) 143: 263–276 DOI 10.1007/s11214-008-9425-1 What Physical Processes Drive the Interstellar Medium in the Local Bubble? D. Breitschwerdt · M.A. de Avillez · B. Fuchs · C. Dettbarn Received: 28 March 2008 / Accepted: 30 July 2008 / Published online: 12 September 2008 © Springer Science+Business Media B.V. 2008 Abstract Recent 3D high-resolution simulations of the interstellar medium in a star form- ing galaxy like the Milky Way show that supernova explosions are the main driver of the structure and evolution of the gas. Its physical state is largely controlled by turbulence due to the high Reynolds numbers of the average flows. For a constant supernova rate a dynam- ical equilibrium is established within 200 Myr of simulation as a consequence of the setup of a galactic fountain. The resulting interstellar medium reveals a typical density/pressure pattern, i.e. distribution of so-called gas phases, on scales of 500–700 pc, with interstellar bubbles being a common phenomenon just like the Local Bubble and the Loop I superbub- ble, which are assumed to be interacting. However, modeling the Local Bubble is special, because it is driven by a moving group, passing through its volume, as it is inferred from the analysis of Hipparcos data. A detailed analysis reveals that between 14 and 19 super- novae have exploded during the last 15 Myr. The age of the Local Bubble is derived from ±0.7 comparison with HI and UV absorption line data to be 14.5 0.4 Myr. We further predict the merging of the two bubbles in about 3 Myr from now, when the interaction shell starts to fragment. The Local Cloud and its companion HI clouds are the consequence of a dynamical instability in the interaction shell between the Local and the Loop I bubble. Keywords Interstellar matter · Superbubbles · Local bubble · Turbulence D. Breitschwerdt () Institut für Astronomie, University of Vienna, Türkenschanzstraße 17, 1180 Vienna, Austria e-mail: [email protected] M.A. de Avillez Department of Mathematics, University of Évora, R. Romão Ramalho 59, 7000 Évora, Portugal B. Fuchs · C. Dettbarn Astronomisches Rechen-Institut, ZAH Universität Heidelberg, Mönchhofstraße 12-14, 69120 Heidelberg, Germany 264 D. Breitschwerdt et al. 1 Introduction The interstellar medium (ISM) in galaxies like our Milky Way is continuously evolv- ing by forming stars in dense and cool clouds, and by enriching the gas with heavy el- ements from stellar explosions thus closing a galactic matter cycle. It has been known for a long time that such feedback mechanisms regulate star formation. More debat- able, and hence less accepted, was the proposed outflow of hot gas from the disk as a galactic fountain (Shapiro and Field 1976;Bregman1980;Kahn1981; de Avillez 2000; Breitschwerdt and Komossa 2000) or even as a wind (e.g. Mathews and Baker 1971; Breitschwerdt et al. 1991; Everett et al. 2008) driven by supernova heated plasma and cos- mic rays. One reason for skepticism lay in the lack of direct observational evidence. Ab- sorption line studies of the galactic halo (e.g. Savage and de Boer 1981) were overwhelm- ingly pointing to negative velocities and hence infalling gas. However, the existence of hot gaseous X-ray halos observed with ROSAT, and now with Chandra and XMM-Newton, ar- gue beyond doubt for outflow in star forming galaxies like NGC 891 (Tüllmann et al. 2006) sometimes labeled as a twin of our galaxy. Why is the outflow from the disk into the halo so important for the evolution of the disk ISM? The simplest answer is: it acts as a pressure release valve. Without it, most of the disk volume, perhaps 50–70% as envisaged by McKee and Ostriker (1977), would be occupied by the hot medium. The evolution of a superbubble like our Local Bubble (LB), in which the solar system is embedded, would be quite different in a high pressure, low density en- vironment. However, the story to tell about the ISM in general and the local ISM (LISM) in particular is not quite as simple and more subtle. In the ISM the estimated Reynolds’ numbers are quite high, about 105–107 (Elmegreen and Scalo 2004), and the gas is highly compressible. Therefore its dynamical and thermal state is largely determined by turbulence, as was first pointed out by von Weizsäcker (1951), which can be both supersonic and super- alfvénic. The implications for the evolution of the ISM, as we shall show in the following sections, are profound. Turbulence does not add only extra pressure to the medium, as it is often simplistically treated, but also transfers energy from larger to smaller scales, or vice versa (inverse cascade), and promotes efficient mixing of chemically enriched material al- lowing for the existence of gas in classical thermally unstable density-temperature regimes. From all we know the LISM on scales of a few hundred parsecs is not vastly different from the general ISM. Star formation seems to be ongoing as can be inferred from the existence of several nearby superbubbles. The Loop I superbubble, with the North Polar Spur being its shock illuminated outer shell and also the largest coherent X-ray structure in the sky, is excited by the Sco Cen association (or one of its distinct subgroups). On the other hand, there is no young stellar cluster known inside the LB. However, a detailed investigation of Hipparcos and radial velocity data of the local stellar population revealed (Berghöfer and Breitschwerdt 2002; Fuchs et al. 2006) the existence of a moving group passing through the volume occupied by the LB during the last 15 million years and generating the LB by the successive explosions of about 20 SNe. This review paper is organized as follows: In Sect. 2 the modeling of the general ISM and the LB are discussed, while in Sect. 3 the results of our models are compared with observations, notably in HI and UV. In Sect. 4 we summarize our results and present some conclusions. 2 Modeling the ISM and the Local Bubble The limitations of analytical and semianalytical models of the state and evolution of the ISM are obvious. One is forced to make assumptions, e.g. of global pressure equilibrium, Physics of the Local Bubble 265 mass balance, etc. like in the so-called three phase model. Transitions from one phase to an- other are instantaneous and regulated by the temperature perturbations at constant pressure of the net heat loss function L = n2 Λ(T ) − nΓ,wheren, T , Λ, Γ are the number density, temperature, interstellar metallicity dependent cooling and heating functions, respectively. ∂L According to Field’s (1965) criterion ( ∂T )P > 0 guarantees stability of a gas phase. How- ever, as we shall see, this is only true if the cooling time is less than the dynamical time scale. In particular turbulence can have a stabilizing effect on gas in the transition regimes. Turbulent flows are fundamentally different in nature from laminar flows. Generally, tur- bulence is fed by the mean flow field via Reynolds stresses, which arise from the randomly fluctuating velocity components in a fluid. Much insight can be gained by looking at the evo- lution and distribution of vorticity (ω =∇×u). In turbulence, the velocities are correlated over a characteristic distance, which can be described by so-called “eddies”. The largest eddies contain the bulk of the kinetic energy and their characteristics are often depending directly on the nature of the source. Since vorticity can neither be created nor destroyed in the interior of a flow, but rather in boundary layers, its evolution is determined by diffusion and stretching of turbulent blobs. The latter leads to the break-up of eddies into smaller structures and a subsequent cascade of energy from the largest to the smallest scales, where viscous dissipation sets in. A simple manifestation of a turbulent flow is shear flow. In the ISM a lot of shear is introduced by interacting flows or the interaction of pressure driven flows with the ambient medium. Therefore a multitude of sources for turbulence exists: stel- lar (e.g. from HII regions, stellar winds, SNe), differential galactic rotation, fluid instabilities (Kelvin-Helmholtz, Parker, magnetorotational instability, etc.), self-gravity, etc. An estimate of their relevance in the ISM has been carried out with the result that SNe dominate by at least a factor of a few (MacLow and Klessen 2004). Therefore it seems justified to start nu- merical investigations of the evolution of the ISM with SNe (and stellar winds) as the driving sources of mass, momentum and energy input. We note, however, that our simulations like those of other groups, suffer from a fundamental shortcoming: we cannot resolve the viscous scale, where energy dissipation occurs. This problem is well known in turbulence studies, and there exist basically two approaches. In direct numerical simulations (DNS), the max- imum Reynolds number is limited to about Re ∼ 1000, since in incompressible turbulence 4/3 4/9 Re ∼ (L/lbox) Nx ,whereL, lbox,andNx are the outer scale at which turbulence is fed in, the box size, and the number of data points required at any instant, respectively. For ISM simulations this is far below realistic Reynolds numbers. In connection with laboratory studies, large eddy simulations (LES) are carried out, in which only the largest scales are resolved, since they only require a tiny fraction of the computing time, and yet carry the most important properties of turbulence, like e.g. the energy flux cascading down to smaller scales. To cut off the flow at some wavenumber kmax can only be done at the price of supply- ing an empirically motivated closure model (from comparison to laboratory studies), which accounts for the hopefully correct energy drainage out of the turbulent cascade.
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