Key words. galaxies: high-redshift – galaxies: starformation – ultraviolet: galaxies – infrared: galaxies – techniques: spectroscopic
In the past decades, great eﬀort has been devoted to the study of the Universe in the redshift range 1 < z < 3. Several studies have shown that this is the epoch when a substantial fraction of galaxy mass assembly took place, and when there is a peak in the evolution of the starformationrate (SFR) density through cos- mic time ( Lilly et al. 1996 ; Madau et al. 1996 ; Dickinson et al. 2003 ; Hopkins & Beacom 2006 ; Daddi et al. 2007 ). In this epoch a critical transformation phase is believed to have occurred, re- vealed by the observed changes in the colour-mass plane where a significant fraction of galaxies moves from the blue cloud of
One of the main objectives in astrophysics during the past 15 years has been to follow the cosmic starformationrate den- sity (SFRD) at ever earlier epochs. But whenever optical data are used, one must apply a dust correction to the luminosity densi- ties (LDs) and a calibration into SFRDs (with their associated uncertainties) to obtain a relevant estimate. Knowing how the dust attenuation evolves in redshift is therefore mandatory if one wishes to study the redshift evolution of the SFRD.
It is interesting to interpret these results in terms of the frequency of nuclear and star-forming activity in galaxies. For this, we consider the distribution of AGN Edding- ton ratios (λǫ∝ ˙ MBH/MBH) and galaxy sSFRs. Recently, Aird et al. (2012) suggested that the λǫ distribution of X-ray AGNs can be described purely as a function of λǫ and redshift; i.e., independently of M∗ (Fig. 3a). 8 This broad distribution for AGNs, which spans over four or- ders of magnitude in λǫ (see also, e.g., Babi´c et al. 2007; Hopkins & Hernquist 2009), contrasts with the distribu- tion of sSFR of galaxies that is remarkably narrow, yet also independent of M∗ (Sargent et al. 2012; Fig. 3a). This is the main reason why the AGN main sequence remains hidden; there are strong changes in s ˙ MBH com- pared to minor changes in the sSFRs of galaxies (Fig. 3). This has the implication that outliers should exist in the MBH-M∗ relation when the SMBH growth has taken advantage over the M∗ growth and vice versa, in qualita-
Subject headings: galaxies: formation — cosmology: observations — infrared: galaxies — galaxies: starburst
— galaxies: evolution
Exploring the relation between the gas content and starformationrate (SFR) of galaxies is crucial to understanding galaxy formation and evolution. This is required to under- stand the nature of SF, the parameters that regulate it, and its possible dependence on local and global galaxy properties (e.g., Silk 1997; Elmegreen 2002; Krumholz & Thompson 2007; McKee & Ostriker 2007). In addition, this information is a critical ingredient of theoretical models of galaxy forma- tion, either based on semianalytical realizations or on numer- ical simulations (e.g., Guiderdoni et al. 1997; Somerville et al. 2001; Monaco et al. 2007; Ocvirk et al. 2008; Dekel et al. 2009; Gnedin et al. 2009; Croton et al. 2006). In fact, the physics associated with the conversion of gas into stars inside galaxies is overwhelmingly complicated, so that theoretical models generally resort to scaling laws that are calibrated us- ing observations of nearby galaxies.
Pappalardo et al.: The stripping age of NGC 4388 11
6.2. Gas dynamics and starformation
With the non-parametric method we find an upper limit for the stripping age of 300 Myr with a smooth decline of starformation during 100-200 Myr. Numerical simulations (Abadi et al. 1999, Vollmer et al. 2001, R¨odiger & Br¨uggen 2006) show that the ISM located at a given distance from the galaxy center is stripped rapidly within a few 10 Myr. Therefore, the smooth decline of the starformationrate derived from the non-parametric method is most probably not physical, but due to regularization of the in- version. A rapid halt of starformation suggested by the dynami- cal models can only be studied with a parametric method, where we make the simplifying assumption that the gas is stripped in- stantaneously. The derived stripping age of 190 Myr is consistent with the extent of the observed H tail (Oosterloo & van Gorkom 2005): a tail extent of 80 kpc in the plane of the sky and along the line of sight with the given stripping age leads to a total velocity
slowly varying, continuous activity that showed a peak about 3 Gyr ago (Cignoni et al. 2006 ), as shown in Fig. 1.1 . One can also refer to the review of Wyse ( 2009 ) for further details. The mechanisms that shape this SFH are still poorly understood today. The regular bursts could be associated with merging events, i.e., the accretion of other smaller galaxies (“dwarf” satellites) on our Milky Way. Not only will these galaxies bring additional stars, they will also briefly destabilize the gas of the Galactic disk and allow it to collapse and form stars more efficiently (this process is discussed further in Section 1.2.2 ). Another explanation which is put forward in Hernandez et al. ( 2000 ) is that, since this SFH is estimated from the solar neighborhood only, i.e., a relatively small region compared to the whole Milky Way, these bursts could correspond to the regular passage of the spiral arms. Somewhat surprisingly, the spiral arms are nothing but density waves inside the disk: they are not representative of the motion of individual stars, but emergent patterns caused by different orbits around the Galactic center (Lindblad 1960 ). When this density wave reaches a given region of the disk, it creates local variations of the gravitational potential and also destabilizes the gas, perhaps in a less efficient way than mergers. This means that, if we were to estimate the SFH from a larger sample of stars not limited to the solar neighborhood (something that Gaïa will soon provide), these variations would vanish, and the SFH of the whole Galaxy would appear much smoother. However, a feature that is expected to remain would be the larger peak that happened 3 Gyr ago. This enhancement of starformation may instead be caused by a major merger, the collision of the Milky Way with another galaxy of similar mass, or by a more intense flow of gas coming from the intergalactic medium (IGM), through a process called “infall” (see, e.g., the discussion in Kennicutt 1983 ). Indeed, our Galaxy must have received large quantities of gas from outside in its past, and probably does so even today. Its present-day starformationrate is currently estimated around SFR = 4 M ⊙ /yr (Diehl et al. 2006 ), while the mass of gas available to form stars is of the order of M gas = 2 × 10 9 M ⊙ (van den Bergh 1999 ). Therefore, at this rate the Milky Way would consume all its gas within 500 Myr (see van den Bergh 1957 , where this problem was first reported). This latter quantity is known as the depletion timescale, tdep. Such short timescale is not specific to the Milky Way: except for a few exceptions which are not representative of star-forming galaxies (e.g., M31 with tdep = 5 Gyr, Pflamm-Altenburg & Kroupa 2009 ), the depletion timescales in the majority of galaxies is typically no more than 1 Gyr (see, e.g., Saintonge et al. 2011 ). This is probably the best evidence that galaxies routinely receive gas from the intergalactic medium.
supernova. On the scale of the galaxy, they found a typical Schmidt-Kennicutt law that only depends on the feedback and not on the implementation of the star-formation process (i.e. density threshold, high-density starformation eﬃciency...). This result suggests that the star-formationrate is regulated only by the feedback processes. Which of these pro- cesses is dominant for the regulation? Fall et al. ( 2010 ) calculated the diﬀerent impact of these feedbacks and found that the radiation pressure is important for massive molecular cloud while at low mass the ionization is dominant. Hopkins et al. ( 2011 ) found that the feedback from stellar winds and supernova is ineﬃcient to regulate starformation because the cooling of the gas is strong enough to evacuate the absorbed energy. On the other hand, the radiation pressure from the absorption of UV/IR photons by the dust and the photo-ionization of the gas are able to disrupt the cloud and prevent the collapse of MCs and GMCs (radiation pressure at high densities for starbursts and galaxies at high redshift and photo-ionization at low densities for the Milky Way and dwarf galaxies, similarly to the results of Fall et al. ( 2010 ). Note that these studies do not take in account the destruc- tion of the dust by the ionization. Therefore, it is still diﬃcult to know if the ionization can be neglected when the radiation pressure seems dominant. A full treatment of the photo-ionization/radiation-pressure physics on a gas/dust mixture is needed to treat that question. Anyway, in the Milky Way, H ii regions are commonly observed in molecular clouds and are probably the dominant process in this range of densities.
Methods. We conducted wide-field mapping of the Aquila, Ophiuchus, and Orion B clouds at ∼0.04 pc resolution in the J = 1−0
transition of HCN, HCO + , and their isotopomers. For each cloud, we derived a reference estimate of the dense gas mass M A V > 8
Herschel , as
well as the strength of the local far-ultraviolet (FUV) radiation field, using Herschel Gould Belt survey data products, and estimated the star-formationrate from direct counting of the number of Spitzer young stellar objects.
5 Institute of Astronomy, The University of Tokyo, 2-21-1 Osawa, Mitaka, Tokyo 181-0015, Japan
Received 26 January 2018 / Accepted 26 March 2018
Variations of starformation activity may happen on a large range of timescales and some of them are expected to be short, that is, a few hundred million years. The study of the physical processes linked to these rapid variations requires large statistical samples to pinpoint galaxies undergoing such transformations. Building upon a previous study, we define a method to blindly identify galaxies that have undergone, and may still be undergoing, a fast downfall of their starformation activity, that is, a more than 80% drop in starformationrate (SFR) occurring in less than 500 Myr. Modeling galaxies’ spectral energy distribution (SED) with a delayed-τ starformation history, with and without allowing an instantaneous SFR drop within the last hundred million years, we isolate 102 candidates out of a subsample of 6680 galaxies classified as “star forming” from the UVJ criterion in the ZFOURGE catalogs. These galaxies are mostly located in the lower part of the SFR-M ∗ main sequence (MS) and extend up to a factor 100 below it. They also lie close to the limit
11 March 2010
The starformation history of nearby early-type galaxies is investigated via numerical modelling. Idealized hydrodynamical N-body simulations with a starformation prescription are used to study the minor merger process (1/10 ≤ M 1 /M 2 ≤ 1/4; M 1 ≤ M 2 ) between a giant galaxy (host) and a less massive spiral galaxy (satellite) with reasonable assumptions for the ages and metallicities of the merger progenitors. We find that the evolution of the starformationrate is extended over several dynamical times and shows peaks which correspond to pericentre passages of the satellite. The newly formed stars are mainly located in the central part of the satellite remnant while the older stars of the initial disk are deposited at larger radii in shell-like structures. After the final plunge of the satellite, starformation in the central part of the rem- nant can continue for several Gyrs depending on the starformation efficiency. Although the mass fraction in new stars is small, we find that the half-mass radius differs from the half-light radius in the V and H bands. Moreover syn- thetic 2D images in J, H, NUV, H β and V bands, using the characteristic filters of the Wide Field Camera 3 (WFC3) on the Hubble Space Telescope (HST), reveal that residual starformation induced by gas-rich minor mergers can be clearly observed during and after the final plunge, especially in the NUV band, for interacting systems at (z ≤ 0.023) over moderate numbers of orbits (∼ 2 orbits correspond to typical exposure times of ∼ 3600 sec). This suggests that WFC3 has the potential to resolve these substructures, characterize plausible past merger episodes, and give clues to the formation of early-type galaxies. Key words: galaxies: formation - galaxies: interactions - galaxies: structure - galaxies: kinematics and dynamics - galaxies: photometry - methods: N-body simulations
In the Magellanic Clouds (MC) the B and Be star populations are not well known or unknown and there is no spectral classification for these objects.
In order to investigate the effects of metallicity, starformation conditions and evolution in B & Be stars, we observed B-type stars, selected on colour- magnitude criteria, with the VLT multifibres GIRAFFE spectrograph. These observations were performed at medium resolution in 2 settings: a “blue” setting centered at 4250 ˚ A to determine the fundamental parameters and a “red” setting centered at 6570 ˚ A for the characterization of the emission of Be stars. Finally, we observed 177 stars in the Large Magellanic Cloud (LMC); among them, 25 are new Be stars, 22 are known Be stars and 121 stars are B stars. In the Small Magellanic Cloud (SMC), we observed 346 stars, of which 90 are new Be stars, 41 are known Be stars and 202 are B stars. In the statistical studies reported in the following sections, we have removed from the samples the binaries (B and Be stars) and we have assumed a random distribution of the inclination angle.
State-of-the art hydrodynamical simulations with improved stellar and AGN feedback can reproduce the cosmic starformation history of the Universe and the galaxy stellar mass func- tion. Hydrodynamic simulations are currently working only above certain mass and spatial resolutions, however, and physical processes on smaller scales are implemented via analytic prescriptions known as ‘sub-grid physics’. EAGLE suite of hydro-simulations with several galaxy formation scenarios empowers us to systematically explore the impact of varying feed- back prescriptions on large representative populations of stellar systems. Strong gravitational lensing is one of the most robust and powerful techniques to measure the total mass and its distribution in galaxies on kpc scales allowing their inner structure and evolution over cosmic time to be studied in detail. I will present and discuss impact of nine different theoretical model in EAGLE (see Mukherjee et al. 2019) on the mass-density and mass-size scaling relation, pos- sible observational systematics (e.g. differences in model-fitting methodologies, differences in filters/bands of the observational surveys, possible lens selection biases, etc.) as well as reso- lution effects in the simulations, that might affect their comparison. We find that models in which stellar feedback becomes inefficient at high gas densities, or weaker AGN feedback with a higher duty cycle, produce strong lenses with total mass density slopes close to isothermal and mass-size relation agreeing with strong lensing observations. I will also demonstrate the differences in comparison of mass-size relation of EAGLE galaxies with strong lensing and non- lensing galaxies from observations. In later part of my talk, I will present our recent analysis of the stellar-to-halomass relation of EAGLE galaxies, which connects the stellar mass M ? of a
Ken Tapping, 2 nd May, 2017
These evenings, after the Sun is gone, you should be able to see a bright, orange-red star low in the west. If the colour is not clear, binoculars will show it well, especially if you throw the image out of focus. That star is in Orion, a winter constellation in the process of vanishing into the Sun’s glare as we get deeper into spring. The star is an old, red giant, popularly called “Betelgeuse”. Astronomer Patrick Moore did not believe any star should be called “Beetle Juice” and consulted some Arab Scholars. The opinion was that “Betelgeuse” came from a French distortion of the Arabic “bit al
That explains the team organisation. The first sub-group is in charge of To support the mission and To change the mode, and the other one is in charge of To estimate the quaternion (Fig. 15).
Figures 15, 16 and 17 are the eFFDB diagrams to illustrate the functional analysis.
As said in section 2.B., the mode of the star tracker described in this paper are the lost-in-space mode and the tracking mode. That is why the function To estimate the quaternion is present in this analysis. There are three functions running at the same time: the estimation of the quaternion from the raw image of the sky, the mission support to enable the algorithm to run, to communicate with the OBDH and to check if the star tracker is all right. The third function is the one that change the current mode when the mission support asks for the change. The scheme 15 explains the functioning of the software when it is in the lost-in-space mode. Three functions are running at the same time. To estimate the quaternion and To support the mission are into an infinite loop, whereas To change the mode occurs only when To support the mission asks for it.
Ken Tapping, 6 th December, 2016
If it happens to be clear, these evenings you can see the Pleiades in the eastern sky. It is a little group of stars, looking like a necklace dropped by a goddess. These stars are often referred to as the “Seven Sisters”, although average eyes may see only six, and good eyes under good conditions will see around ten. Binoculars will show the group to have hundreds of members. The Pleiades is an example of an “open star cluster” – open because its members are spread out and easy to see. The stars in the Pleiades are “siblings”, born from the same cloud of gas and dust, which collapsed about 100 million years ago. As it collapsed it broke into hundreds of little “cloudlets” each of which eventually formed stars. The birth place is now long gone, but we think it looked very much like the Orion Nebula, a huge cloud of glowing gas where we can see almost all stages in star
Accepted 2013 April 8. Received 2013 March 26; in original form 2012 October 9
A B S T R A C T
We explore the effects of dynamical evolution in dense clusters on the companion mass ratio distribution (CMRD) of binary stars. Binary systems are destroyed by interactions with other stars in the cluster, lowering the total binary fraction and significantly altering the initial semimajor axis distribution. However, the shape of the CMRD is unaffected by dynamics; an equal number of systems with high mass ratios are destroyed compared to systems with low mass ratios. We might expect a weak dependence of the survivability of a binary on its mass ratio because its binding energy is proportional to both the primary and secondary mass components of the system. However, binaries are broken up by interactions in which the perturbing star has a significantly higher energy (by a factor of 10, depending on the particular binary properties) than the binding energy of the binary, or through multiple interactions in the cluster. We therefore suggest that the shape of the observed binary CMRD is an outcome of the starformation process and should be measured in preference to the distributions of orbital parameters, such as the semimajor axis distribution.
4.1. StarFormation History via CMD-fitting
Here, we apply the CMD-fitting package StarFISH (Harris & Zaritsky 2001 ) to our Leo IV photometry within the half- light radius to determine its SFH and metallicity evolution. As discussed in previous works, StarFISH uses theoretical isochrones (we use those of Girardi et al. 2004 , although any may be used) to construct artificial CMDs with different combinations of distance, age, and metallicity. Once convolved with the observed photometric errors and completeness (using the artificial star tests of Section 2 ), these theoretical CMDs are converted into realistic model CMDs which can be directly compared to the data on a pixel-to-pixel basis, after binning each into Hess diagrams. We use the Poisson statistic of Dolphin ( 2002 ) as our fit statistic. The best-fitting linear combination of model CMDs is determined through a modification of the standard AMOEBA algorithm (Press et al. 1997 ; Harris & Zaritsky 2001 ). Several steps are taken to determine the uncertainties in StarFISH fits, which are discussed in detail in previous work—see Harris & Zaritsky ( 2009 ); de Jong et al.
cores ( André et al. 2010 ). While molecular clouds have been known to exhibit large-scale filamentary structures for quite some time (e.g. Schneider & Elmegreen 1979 ; Myers 2009 , and references therein), Herschel observations now demonstrate that these filaments are truly ubiquitous in the cold ISM (e.g. Molinari et al. 2010 ; Henning et al. 2010 ; Hill et al. 2011 ), prob- ably make up a dominant fraction of the dense gas in molec- ular clouds (e.g. Schisano et al. 2014 ; Könyves et al. 2015 ), and present a high degree of universality in their properties (e.g. Arzoumanian et al. 2011 ). This means that interstellar fil- aments probably play a central role in the starformation pro- cess (e.g. André et al. 2014 ). A detailed analysis of their ra- dial column density profiles shows that, at least in the nearby clouds of the Gould Belt, filaments are characterized by a very narrow distribution of inner widths W with a typical FWHM value ∼0.1 pc (much higher than the ∼0.01 pc resolution pro- vided by Herschel at the distance ∼140 pc of the nearest clouds) and a dispersion lower than a factor of 2 ( Arzoumanian et al. 2011 ; Koch & Rosolowsky 2015 ). The origin of this common inner width of interstellar filaments is not yet well understood. A possible interpretation is that it corresponds to the sonic scale below which interstellar turbulence becomes subsonic in dif- fuse, non-star-forming molecular gas (cf. Padoan et al. 2001 ; Federrath 2016 ). Alternatively, this characteristic inner width of filaments may be set by the dissipation mechanism of magneto- hydrodynamic (MHD) waves (e.g. Hennebelle & André 2013 ). A possible manifestation of such MHD waves may have been found in the form of braided velocity-coherent substructure in the Taurus B211–3 filament ( Hacar et al. 2013 ). Another major result from Herschel in nearby clouds is that most (>75%) low- mass prestellar cores and protostars are found in dense, super- critical filaments for which the mass per unit length M line ex-
The impact of massive stars on their environments can be ei- ther destructive or constructive (Gorti & Hollenbach 2002; Tan & McKee 2001). If constructive, a massive star will favour subsequent starformation via energetic phenomena such as winds and radiation, leading to a local increase in pressure. Indeed, many physical processes can trigger starformation (see Elmegreen 1998 for a review). The collect and collapse process is one of them. In this process (first proposed by Elmegreen & Lada 1977), the expansion of an H region generates the formation of a layer of gas and dust that is accumulated be- tween the ionization front (IF) and the shock front (SF) that precedes the IF on the neutral side. This compressed layer may become gravitationally unstable along its surface and then frag- ment. This process is interesting as it allows the formation of massive fragments out of a previously uniform medium, thus
– In extremely active environments with normal metal-
licity, the aromatic band carriers can also be destroyed or experience chemical transformations, but quantita- tive estimations of these effects are yet unavailable. However, this relationship between mid-IR fluxes and SFRs that strictly holds only in normal disks can be useful to interpret surveys made in the two filters LW3 (15 µm) and LW2 (7 µm). For galaxies at high redshifts (z ' 1.2), the LW2 rest frame emission is shifted to the LW3 band- pass. Hence, the calibration given here must provide a lower limit for the true SFR since, for galaxies with greater starformation activity than in the present sample, the en- ergy redistribution favors the LW3 band, as more energy is reradiated by VSGs. Indeed, Boulanger et al. (1998b) have presented observations in resolved Galactic regions (a dif- fuse cloud and four photodissociation regions) which show that the emission in UIBs tends to rise with the ultraviolet energy density, linearly at low energy densities and more slowly at higher values. The threshold for this transition,