The SILCC project (SImulating the Life-Cycle of molecular Clouds) aims at a more self-consistent understanding of the interstellar medium (ISM) on small scales and its link to galaxy evolution. We present three-dimensional (magneto)hydrodynamic simulations of the ISM in a vertically stratified box including self-gravity, an external potential due to the stellar component of the galactic disc, and stellar feedback in the form of an interstellar radiation field and supernovae (SNe). The cooling of the gas is based on a chemical network that follows the abundances of H + , H, H 2 , C + , and CO and takes shielding into account consistently. We vary the SN feedback by comparing different SN rates, clustering and different positioning, in particular SNe in density peaks and at random positions, which has a major impact on the dynamics. Only for random SN positions the energy is injected in sufficiently low-density environments to reduce energy losses and enhance the effective kinetic coupling of the SNe with the gas. This leads to more realistic velocity dispersions (σ HI ≈ 0.8σ 300−8000 K ∼ 10−20 km s −1 , σ Hα ≈ 0.6σ 8000−3×10 5 K ∼ 20 − 30 km s −1 ), and strong outflows with mass loading factors (ratio of outflow to star formation rate) of up to 10 even for solar neighbourhood conditions. Clustered SNe abet the onset of outflows compared to individual SNe but do not influence the net outflow rate. The outflows do not contain any molecular gas and are mainly composed of atomic hydrogen. The bulk of the outflowing mass is dense (ρ ∼ 10 −25 − 10 −24 g cm −3 ) and slow (v ∼ 20 − 40 km s −1 ) but there is a high-velocity tail of up to v ∼ 500 km s −1 with ρ ∼ 10 −28 − 10 −27 g cm −3 .

We present a hydrodynamical simulation of the turbulent, magnetized, supernova (SN)-driven interstellar medium (ISM) in a stratified box that dynamically couples the injection and evolution of cosmic rays (CRs) and a self-consistent evolution of the chemical composition. CRs are treated as a relativistic fluid in the advection-diffusion approximation. The thermodynamic evolution of the gas is computed using a chemical network that follows the abundances of H + , H, H 2 , CO, C + , and free electrons and includes (self-)shielding of the gas and dust. We find that CRs perceptibly thicken the disk with the heights of 90% (70%) enclosed mass reaching 1.5 kpc ( 0.2 kpc). The simulations indicate that CRs alone can launch and sustain strong outflows of atomic and ionized gas with mass loading factors of order unity, even in solar neighborhood conditions and with a CR energy injection per SN of 10 50 erg, 10% of the fiducial thermal energy of an SN. The CR-driven outflows have moderate launching velocities close to the midplane ( 100 km s −1 ) and are denser (ρ ∼ 10 −24 − 10 −26 g cm −3 ), smoother, and colder than the (thermal) SN-driven winds. The simulations support the importance of CRs for setting the vertical structure of the disk as well as the driving of winds.

Feedback from massive stars is believed to be a key element in the evolution of molecular clouds. We use high‐resolution 3D smoothed particle hydrodynamics simulations to explore the dynamical effects of a single O7 star‐emitting ionizing photons at 1049 s−1 and located at the centre of a molecular cloud with mass 104 M⊙ and radius 6.4 pc; we also perform comparison simulations in which the ionizing star is removed. The initial internal structure of the cloud is characterized by its fractal dimension, which we vary between D=2.0 and 2.8, and the standard deviation of the approximately log‐normal initial density PDF, which is σ10 = 0.38 for all clouds. (i) As regards star formation, in the short term ionizing feedback is positive, in the sense that star formation occurs much more quickly (than in the comparison simulations), in gas that is compressed by the high pressure of the ionized gas. However, in the long term ionizing feedback is negative, in the sense that most of the cloud is dispersed with an outflow rate of up to ∼10−2 M⊙yr−1, on a time‐scale comparable with the sound‐crossing time for the ionized gas (∼1−2 Myr ), and triggered star formation is therefore limited to a few per cent of the cloud's mass. We will describe in greater detail the statistics of the triggered star formation in a companion paper. (ii) As regards the morphology of the ionization fronts (IFs) bounding the H ii region and the systematics of outflowing gas, we distinguish two regimes. For low D≲2.2, the initial cloud is dominated by large‐scale structures, so the neutral gas tends to be swept up into a few extended coherent shells, and the ionized gas blows out through a few large holes between these shells; we term these H ii regions shell dominated. Conversely, for high D≳2.6, the initial cloud is dominated by small‐scale structures, and these are quickly overrun by the advancing IF, thereby producing neutral pillars protruding into the H ii region, whilst the ionized gas blows out through a large number of small holes between the pillars; we term these H ii regions pillar dominated. (iii) As regards the injection of bulk kinetic energy, by ∼1 Myr, the expansion of the H ii region has delivered a mass‐weighted rms velocity of ∼6 km s−1; this represents less than 0.1 per cent of the total energy radiated by the O7 star.

We study the impact of stellar winds and supernovae on the multi-phase interstellar medium using three-dimensional hydrodynamical simulations carried out with FLASH. The selected galactic disc region has a size of (500 pc) 2 × ±5 kpc and a gas surface density of 10 M pc −2 . The simulations include an external stellar potential and gas self-gravity, radiative cooling and diffuse heating, sink particles representing star clusters, stellar winds from these clusters which combine the winds from individual massive stars by following their evolution tracks, and subsequent supernova explosions. Dust and gas (self-)shielding is followed to compute the chemical state of the gas with a chemical network. We find that stellar winds can regulate star (cluster) formation. Since the winds suppress the accretion of fresh gas soon after the cluster has formed, they lead to clusters which have lower average masses (10 2 −10 4.3 M ) and form on shorter timescales (10 −3 −10 Myr). In particular we find an anti-correlation of cluster mass and accretion time scale. Without winds the star clusters easily grow to larger masses for ∼ 5 Myr until the first supernova explodes. Overall the most massive stars provide the most wind energy input, while objects beginning their evolution as B type stars contribute most of the supernova energy input. A significant outflow from the disk (mass loading 1 at 1 kpc) can be launched by thermal gas pressure if more than 50% of the volume near the disc mid-plane can be heated to T > 3 × 10 5 K. Stellar winds alone cannot create a hot volume-filling phase. The models which are in best agreement with observed star formation rates drive either no outflows or weak outflows.

The SILCC project (SImulating the Life-Cycle of molecular Clouds) aims at a more selfconsistent understanding of the interstellar medium (ISM) on small scales and its link to galaxy evolution. We simulate the evolution of the multi-phase ISM in a (500 pc) 2 × ± 5 kpc region of a galactic disc, with a gas surface density of Σ GAS = 10 M ⊙ /pc 2 . The Flash 4.1 simulations include an external potential, self-gravity, magnetic fields, heating and radiative cooling, time-dependent chemistry of H 2 and CO considering (self-) shielding, and supernova (SN) feedback. We explore SN explosions at different (fixed) rates in high-density regions (peak), in random locations (random), in a combination of both (mixed), or clustered in space and time (clustered). Only random or clustered models with self-gravity (which evolve similarly) are in agreement with observations. Molecular hydrogen forms in dense filaments and clumps and contributes 20% -40% to the total mass, whereas most of the mass (55% -75%) is in atomic hydrogen. The ionised gas contributes <10%. For high SN rates (0.5 dex above Kennicutt-Schmidt) as well as for peak and mixed driving the formation of H 2 is strongly suppressed. Also without self-gravity the H 2 fraction is significantly lower (∼ 5%). Most of the volume is filled with hot gas (∼90% within ±2 kpc). Only for random or clustered driving, a vertically expanding warm component of atomic hydrogen indicates a fountain flow. Magnetic fields have little impact on the final disc structure. However, they affect dense gas (n 10 cm −3 ) and delay H 2 formation. We highlight that individual chemical species, in particular atomic hydrogen, populate different ISM phases and cannot be accurately accounted for by simple temperature-/density-based phase cut-offs.

We present 3D "zoom-in" simulations of the formation of two molecular clouds out of the galactic interstellar medium. We model the clouds -identified from the SILCC simulations -with a resolution of up to 0.06 pc using adaptive mesh refinement in combination with a chemical network to follow heating, cooling, and the formation of H 2 and CO including (self-) shielding. The two clouds are assembled within a few million years with mass growth rates of up to ∼ 10 −2 M yr −1 and final masses of ∼ 50 000 M . A spatial resolution of 0.1 pc is required for convergence with respect to the mass, velocity dispersion, and chemical abundances of the clouds, although these properties also depend on the cloud definition such as based on density thresholds, H 2 or CO mass fraction. To avoid grid artefacts, the progressive increase of resolution has to occur within the free-fall time of the densest structures (1 -1.5 Myr) and 200 time steps should be spent on each refinement level before the resolution is progressively increased further. This avoids the formation of spurious, large-scale, rotating clumps from unresolved turbulent flows. While CO is a good tracer for the evolution of dense gas with number densities n 300 cm −3 , H 2 is also found for n 30 cm −3 due to turbulent mixing and becomes dominant at column densities around 30 -50 M pc −2 . The CO-to-H 2 ratio steadily increases within the first 2 Myr whereas X CO 1 -4 × 10 20 cm −2 (K km s −1 ) −1 is approximately constant since the CO(1-0) line quickly becomes optically thick.

We present simulations of stable isothermal clouds exposed to ionizing radiation from a discrete external source, and identify the conditions that lead to Radiatively Driven Implosion and Star Formation. We use the Smoothed Particle Hydrodynamics code SEREN (Hubber et al. 2010) and the HEALPix-based photoionization algorithm described in Bisbas et al. (2009). We find that the incident ionizing flux is the critical parameter determining the evolution; high fluxes disperse the cloud, whereas low fluxes trigger star formation. We find a clear connection between the intensity of the incident flux and the parameters of star formation.

We use hydrodynamical simulations in a (256 pc) 3 periodic box to model the impact of supernova (SN) explosions on the multi-phase interstellar medium (ISM) for initial densities n = 0.5-30 cm −3 and SN rates 1-720 Myr −1 . We include radiative cooling, diffuse heating, and the formation of molecular gas using a chemical network. The SNe explode either at random positions, at density peaks, or both. We further present a model combining thermal energy for resolved and momentum input for unresolved SNe. Random driving at high SN rates results in hot gas (T 10 6 K) filling > 90% of the volume. This gas reaches high pressures (10 4 < P/k B < 10 7 K cm −3 ) due to the combination of SN explosions in the hot, low density medium and confinement in the periodic box. These pressures move the gas from a two-phase equilibrium to the single-phase, cold branch of the cooling curve. The molecular hydrogen dominates the mass (> 50%), residing in small, dense clumps. Such a model might resemble the dense ISM in high-redshift galaxies. Peak driving results in huge radiative losses, producing a filamentary ISM with virtually no hot gas, and a small molecular hydrogen mass fraction ( 1%). Varying the ratio of peak to random SNe yields ISM properties in between the two extremes, with a sharp transition for equal contributions. The velocity dispersion in Hi remains 10 km s −1 in all cases. For peak driving the velocity dispersion in H α can be as high as 70 km s −1 due to the contribution from young, embedded SN remnants.