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Characterizing the [line as a narrow plus a broad coniponen, Dwasawaοal.(1996) fouud tha the intensity of the narrow component correlates with the coutimmun flux. whereas the intensity of the broad conmonent possibly auti-correlates with the continuum flux. | Characterizing the line as a narrow plus a broad component, \citet[]{iwa96}
found that the intensity of the narrow component correlates with the continuum flux, whereas the intensity of the broad component possibly anti-correlates with the continuum flux. |
This «observational result is also qualitatively consistent with the out-Howiue magnetic flare model. | This observational result is also qualitatively consistent with the out-flowing magnetic flare model. |
The A-rav οwission dn NGCSSLS is also required | The X-ray emission in NGC5548 is also required — |
For an ensemble of relativistic electrons wilh a power-law distribution of energy NCL) and an isotropic velocity distribution. their svnchrotron emission coellicient is where @ is the angle between (he magnetic field /3 and the line of sight to the observer. | For an ensemble of relativistic electrons with a power-law distribution of energy $N(E)=N_\circ E^{-\gamma}$ and an isotropic velocity distribution, their synchrotron emission coefficient is where $\theta$ is the angle between the magnetic field $B$ and the line of sight to the observer. |
The resulting power-law spectrum has a lorm 5,xva" where a=(4—1)/2 is the spectral index. | The resulting power-law spectrum has a form $S_\nu \propto \nu^{-\alpha}$ where $\alpha =
(\gamma -1)/2$ is the spectral index. |
proposed the acceleration of cosmic ravs by supersonic shocks (of velocity ος) as (he main mechanism lor non-thermal svnchrotron emission [rom supernova renmants ancl derived (he volume enmissivilv as. Assuming supernova remnants dominate the non-thermal svnehrotron emission from ealaxies. derived a relation μηxvy, where vsy is the Type II supernova rale per vear (see his Eq. | proposed the acceleration of cosmic rays by supersonic shocks (of velocity $v_s$ ) as the main mechanism for non-thermal synchrotron emission from supernova remnants and derived the volume emissivity as, Assuming supernova remnants dominate the non-thermal synchrotron emission from galaxies, derived a relation $S_{nth}\propto \nu_{SN}$ where $\nu_{SN}$ is the Type II supernova rate per year (see his Eq. |
17). | 17). |
llowever. this relation predicts a vy, more than an order of magnitude too large for our Galaxy. and Condon. proposed. a revised relation based on Galactic normalization (ie. his Eq. | However, this relation predicts a $\nu_{SN}$ more than an order of magnitude too large for our Galaxy, and Condon proposed a revised relation based on Galactic normalization (i.e. his Eq. |
13). which can be re-written as where D, is luminosity distance in Alpe. | 18), which can be re-written as where $D_L$ is luminosity distance in Mpc. |
The first ancl the foremost important step in all photometric redshift technique is establishing a robust and reliable template. | The first and the foremost important step in all photometric redshift technique is establishing a robust and reliable template. |
ev SED template parameters for the starburst SED template are Z; and 3 for the dust emission (eq. 4)). | Key SED template parameters for the starburst SED template are $T_d$ and $\beta$ for the dust emission (Eq. \ref{eq:Sd3}) ), |
Z aud £V for thermal Iree-Iree enussion (Eq. 7)). | $T_e$ and $EM$ for thermal free-free emission (Eq. \ref{eq:frfr2}) ), |
and rss and α for (he non-thermal svnchrotron emission (Eq. 10)). | and $\nu_{SN}$ and $\alpha$ for the non-thermal synchrotron emission (Eq. \ref{eq:nth1}) ). |
These parameters and (heir dispersions are determined by deriving best fit SED models for the 23B Ih selected starburst galaxies whose subnun/FIR data are readily available in (he literature. | These parameters and their dispersions are determined by deriving best fit SED models for the 23 IR selected starburst galaxies whose submm/FIR data are readily available in the literature. |
Flux density for dust emission Sy can be derived directly using Eq. | Flux density for dust emission $S_d$ can be derived directly using Eq. |
3. bv specilving the source size Q4. dust temperature Z;. and enissivitv ο, | \ref{eq:Sd2} by specifying the source size $\Omega_d$, dust temperature $T_d$, and emissivity $\beta$. |
The effective source solid angle ο for a given source is derived from the observed. 100 jan flux density. assuming dust emission is nearly optically thick (see 2.1)). | The effective source solid angle $\Omega_d$ for a given source is derived from the observed 100 $\mu$ m flux density, assuming dust emission is nearly optically thick (see \ref{sec:Sdust}) ). |
For example. the prototvpical ultraluminous galaxy | For example, the prototypical ultraluminous galaxy |
mass.my. lo the projectile mass. M. are myi/M21. | mass,$m_e$, to the projectile mass, $M$ , are $m_e/M \gg$ 1. |
Material ejected by impacts max explain observed activitv in some asteroids. | Material ejected by impacts may explain observed activity in some asteroids. |
In (his section we aim to obtain a relation between the inpactor properties ancl the resulting brightening caused by ejected material. | In this section we aim to obtain a relation between the impactor properties and the resulting brightening caused by ejected material. |
In an impact. the bulk of the ejecta travel at the lowest speeds. | In an impact, the bulk of the ejecta travel at the lowest speeds. |
For equal target. aud projectile densities. (he mass of ejecta. my. traveling faster than a given speed. ο. can be roughly expressed bv a power law in which A and U are the impactor mass aud speed and sl 0.01 is a constant (Llousen and Holsapple 2011). | For equal target and projectile densities, the mass of ejecta, $m_e$, traveling faster than a given speed, $v$, can be roughly expressed by a power law in which $M$ and $U$ are the impactor mass and speed and $A \sim$ 0.01 is a constant (Housen and Holsapple 2011). |
The index a depends «Πρίν on the properties of the target but. is reasonably well approximated for a range of materials as a = -1.5. | The index $\alpha$ depends slightly on the properties of the target but is reasonably well approximated for a range of materials as $\alpha$ = -1.5. |
Only ejecta traveling with ο can escape [rom an asteroid to produce an increase in the scattering cross-section. while (he rest must fall back to coat the surface around the impact site. | Only ejecta traveling with $v \ge v_e$ can escape from an asteroid to produce an increase in the scattering cross-section, while the rest must fall back to coat the surface around the impact site. |
With the escape velocity given by ib is clear (hat Equation (5)) defines. for a given target density. p and impact speed CU. a relation between the impact vield and (he radius. r. of the impacted asteroid. | With the escape velocity given by it is clear that Equation \ref{housen}) ) defines, for a given target density $\rho$ and impact speed $U$, a relation between the impact yield and the radius, $r$, of the impacted asteroid. |
A relation between m. and the scattering cross-section of the ejecta and hence the change in brightness caused by mipact. can be established. | A relation between $m_e$ and the scattering cross-section of the ejecta and hence the change in brightness caused by impact, can be established. |
The size distribution of the ejecta from NASA's Deep Impact mission to comet 9P/Tempel 1 has been modeled as apower law (Ixadono et al. | The size distribution of the ejecta from NASA's Deep Impact mission to comet 9P/Tempel 1 has been modeled as apower law (Kadono et al. |
2010). | 2010). |
Their Figure 4 gives a differential power law size inclex ¢ = 3.7 over the 1 σα< 100 jm size range. while other investigators have found slightly steeper distributions (Lisse et al. | Their Figure 4 gives a differential power law size index $q$ = 3.7 over the 1 $\le a \le$ 100 $\mu$ m size range, while other investigators have found slightly steeper distributions (Lisse et al. |
2006). | 2006). |
Separately. the size distribution of the ejecta from P/2010 A2 has been modeled as a power law with q = 3.32:0.2 for 1<a<10 mm (Jewitt et al. | Separately, the size distribution of the ejecta from P/2010 A2 has been modeled as a power law with $q$ = $\pm$ 0.2 for $1 \le a \le 10$ mm (Jewitt et al. |
2010). | 2010). |
A wider range of power laws (from q = 3 to as steep as q = 6) has been reported in small-scale. hypervelocity. impact. experiments (Takasawa et al. | A wider range of power laws (from $q$ = 3 to as steep as $q$ = 6) has been reported in small-scale, hypervelocity impact experiments (Takasawa et al. |
2011). | 2011). |
For the sake of the present discussion. we adopt ο = 3.5 ancl find that the relation between the scattering cross-section. C. and themass of particles having radii in the range dyjy<a<(54, 1s where —7=((0,,:,0,,,,)E> 7. | For the sake of the present discussion, we adopt $q$ = 3.5 and find that the relation between the scattering cross-section, $C_e$ and themass of particles having radii in the range $a_{min} \le a \le a_{max}$ is where $\overline{a} = (a_{min} a_{max})^{1/2}$ . |
For example. with. ey), = 0.1 jm. tye» = 0.1 m. —qd — m. or | For example, with $a_{min}$ = 0.1 $\mu$ m, $a_{max}$ = 0.1 m, $\overline{a}$ = $^{-4}$ m, or |
'Table3 give / ancl corresponding values of £440) ancl Q[4 derived. from. Hamuy's (2001) Figures 5.7 ancl 5.8 except 1999e7 for which the values were calculated from the data of Leonard οἱ al. ( | $\,$ 3 give $t$ and corresponding values of $F_{41}(t)$ and $Q$ derived from Hamuy's (2001) Figures 5.7 and 5.8 except $\,$ 1999gi for which the values were calculated from the data of Leonard et al. ( |
20025). | 2002b). |
If the value of Marin was known. one could casily [ind the distance D from (7)). | If the value of ${\cal M}_{\mathrm{Ni0}}$ was known, one could easily find the distance $D$ from $\,$ \ref{MNi}) ). |
So. we have to look for à way to estimate Mopio independently. | So, we have to look for a way to estimate ${\cal M}_{\mathrm{Ni0}}$ independently. |
Lt sccms reasonable to assume that the supernova explosion energy { should correlate with Ανω produced during the explosion. | It seems reasonable to assume that the supernova explosion energy $E$ should correlate with ${\cal M}_{\mathrm{Ni0}}$ produced during the explosion. |
This means that where f represents rather than a strict mathematical relation. | This means that where $f$ represents rather than a strict mathematical relation. |
Inserting this expression for. ££ into (11) and. using (4)) for AA. we obtain an equation which can be solved for D when Vo uly. tou. Af. and Q are known from observations. | Inserting this expression for $E$ into $\,$ \ref{Evtu}) ) and using $\,$ \ref{MVAD}) ) for $M_V$, we obtain an equation which can be solved for $D$ when $V-A_V$ , $u_{\mathrm{ph}}$ , $\Delta t$, and $Q$ are known from observations. |
Then for given D. we can find ££. M. H. and Maio from Eqs.(1)) (3)). (43). and. (7)). respectively. | Then for given $D$, we can find $E$, ${\cal M}$, $R$, and ${\cal M}_{\mathrm{Ni0}}$ from $\,$ \ref{Evtu}) \ref{Rvtu}) ), $\,$ \ref{MVAD}) ), and $\,$ \ref{MNi}) ), respectively. |
What can be said about the function. f(Mio) at oesent. when the details of the IE mechanism remain still ambiguous? | What can be said about the function $f\left({\cal M}_{\mathrm{Ni0}}\right)$ at present, when the details of the $\,$ II mechanism remain still ambiguous? |
First of all. it. is reasonable to assume hat a good fraction of /£ comes from the recombination of [ree neutrons and. Nijustatthebottomprotonsinto of the envelope to be finally expelled (Nadyozhin. 1978. Bethe 1996). | First of all, it is reasonable to assume that a good fraction of $E$ comes from the recombination of free neutrons and protons into } just at the bottom of the envelope to be finally expelled (Nadyozhin 1978, Bethe 1996). |
The hydrodynamical modelling of the collapse (Nadvozhin. LOTS) have indicated that under. favourable conditions a neutron-proton shell could be accumulated just under the steady. accreting shock wave. | The hydrodynamical modelling of the collapse (Nadyozhin 1978) have indicated that under favourable conditions a neutron-proton shell could be accumulated just under the steady accreting shock wave. |
When the mass of such a shell reaches some critical value (presumably of the order of 0.1. )) the shell can become unstable in respect to recombining into the "iron group elements (specifically intoNi) to supply the stalled shock wave with the energy of =10" erg necessary to trigger the supernova. | When the mass of such a shell reaches some critical value (presumably of the order of $\approx$ ) the shell can become unstable in respect to recombining into the ”iron group" elements (specifically into }) to supply the stalled shock wave with the energy of $\approx 10^{51}$ erg necessary to trigger the supernova. |
Here. there is a physical analoew with the origin of planctary nebulae from red. giants where the energy. from the recombination of hydrogen and helium causes the expulsion of a red giant rarefied. envelope. | Here, there is a physical analogy with the origin of planetary nebulae from red giants where the energy from the recombination of hydrogen and helium causes the expulsion of a red giant rarefied envelope. |
“Phe recent study. (Imshennik 2002. and references therein) of the “neutrino crown” the region enclosed within neutrinosphere and accreting shock. turns out to be in line with such a picture of the supernova mechanism. | The recent study (Imshennik 2002, and references therein) of the ”neutrino crown" – the region enclosed within neutrinosphere and accreting shock, turns out to be in line with such a picture of the supernova mechanism. |
However. some Ni can be produced through the explosive carbon-oxvecn burning induced. by the outgoing shock wave. | However, some Ni can be produced through the explosive carbon-oxygen burning induced by the outgoing shock wave. |
In this case the energy release per unit Ni mass is lower by an order of magnitude than for the neutron-proton recombination. | In this case the energy release per unit Ni mass is lower by an order of magnitude than for the neutron-proton recombination. |
The energy released neutron-proton α "NIbyουof mass Main. is given by ‘Thus. the ooducetion. of only ~ "Ni 0.06.A4;of is sullicient. to. provide the standard. explosion energy. of 107 erg. | The energy released by the neutron-proton recombination, producing a } mass of ${\cal M}_{\mathrm{Ni0}}$, is given by Thus, the production of only $\sim\NMS{0.06}$ of } is sufficient to provide the standard explosion energy of $10^{51}\,$ erg. |
The current ivdrodyvnamic models of the LL explosions (Woosley Weaver 1995: Rauscher ct al. | The current hydrodynamic models of the $\,$ II explosions (Woosley Weaver 1995; Rauscher et al. |
2002) do not show a correlation between A anc Mio because in these "Nicomesfronexplosivemodels silicon ancl carbon-oxvecn burning near to the envelope. bottom and its vield is sensitive to the mass cut. point. | 2002) do not show a correlation between $E$ and ${\cal M}_{\mathrm{Ni0}}$ because in these models } comes from explosive silicon and carbon-oxygen burning near to the envelope bottom and its yield is sensitive to the mass cut point. |
The photometrical anc spectroscopical properties of the SN models are virtually indepencent of the mass cut. | The photometrical and spectroscopical properties of the SN models are virtually independent of the mass cut. |
On the contrary. the nucleosvnthesis vields are very sensitive to the mass cut. | On the contrary, the nucleosynthesis yields are very sensitive to the mass cut. |
In the current SN models the explosion is usually simulated. by locating a piston attheinternal boundary m= Mau. | In the current SN models the explosion is usually simulated by locating a piston attheinternal boundary $m={\cal M}_{\mathrm {cut}}$ . |
The piston moves with time according to a | The piston moves with time according to a |
The cold dark matter (CDM). cosmogony predicts. that cosmic structures form. hicrarchically through a succession of mergers anc acerctions. | The cold dark matter (CDM) cosmogony predicts that cosmic structures form hierarchically through a succession of mergers and accretions. |
“The mass of the Milkv Way's own dark halo is currently constrained to ~15107AL. (e.g. 2)). | The mass of the Milky Way's own dark halo is currently constrained to $\sim1-5\times10^{12}\,M_{\sun}$ (e.g. \citealt{Guo09arxiv}) ). |
Numerical simulations predict a wealth of dark matter substructures (self-bound subhaloes) surviving in haloes of this mass (e.g. 2: ???)). | Numerical simulations predict a wealth of dark matter substructures (self-bound subhaloes) surviving in haloes of this mass (e.g. \citealt{Gao2004a}; \citealt{Gao2004b, DKM08Nature,
volker08Aq}) ). |
These substructures have a power-law mass function: scaling this to a satellite galaxy Iuminosity function (LE) by adopting a fixed. mass-to-light ratio consistent with the brightest objects overpredicts the number of satellite galaxies of the Milkv Way (MW) and AI31 by a factor of several hundred: —- the so-called "missing satellite’ problem (e.g. ο τι for à recent review see 2). | These substructures have a power-law mass function; scaling this to a satellite galaxy luminosity function (LF) by adopting a fixed mass-to-light ratio consistent with the brightest objects overpredicts the number of satellite galaxies of the Milky Way (MW) and M31 by a factor of several hundred $-$ the so-called `missing satellite' problem (e.g. \citealt{Klypin1999apj,
Moore1999apj}; ; for a recent review see \citealt{Kravtsov2010Review}) ). |
For many vears. theoretical models of galaxy formation have predicted. that the star formation elliciency in. low-mass haloes can be stronely suppressed through a combination of photoionization (?)) and supernova feedback. (2)). | For many years, theoretical models of galaxy formation have predicted that the star formation efficiency in low-mass haloes can be strongly suppressed through a combination of photoionization \citealt{Efstathiou1992}) ) and supernova feedback \citealt{WhiteRees1978}) ). |
"Ehis suppression renders many such haloes permanently dark and may solve the apparent discrepancy (c.g. οτοτὸτι Tu 7?)) | This suppression renders many such haloes permanently `dark' and may solve the apparent discrepancy (e.g. \citealt{Kauffmann1993, Bullock2000, Gnedin2000, Benson02sats,
Okamoto09Frenk, Maccio2009ExplainLF}; \citealt{LiHelmiLucia2010}; \citealt{Stringer09arxiv}) ). |
Recent discoveries of low luminosity satellites in he MW and. M31 ancl corrections for incompleteness (e.g. ?7)) imply reasonable agreement between observations and heoretical predictions for AIW-like haloes (2)). | Recent discoveries of low luminosity satellites in the MW and M31 and corrections for incompleteness (e.g. \citealt{Koposov2008MWLF, Tollerud2008}) ) imply reasonable agreement between observations and theoretical predictions for MW-like haloes \citealt{Benson02sats}) ). |
Llowever. cause they rely on uncertain barvonie physics. and currently limited data. these comparisons so far only provide an indirect ancl uncertain test of the abundance and mass unction of CDM substructures. | However, because they rely on uncertain baryonic physics and currently limited data, these comparisons so far only provide an indirect and uncertain test of the abundance and mass function of CDM substructures. |
Since gravitational lensing races mass. regardless of its luminositv. it can provide a xowerful alternative tool to investigate otherwise "hidden subhaloes and clirectly test the CDM moclel (e.g. ο??)) | Since gravitational lensing traces mass, regardless of its luminosity, it can provide a powerful alternative tool to investigate otherwise `hidden' subhaloes and directly test the CDM model (e.g. \citealt{MKA09SubImageAnomaly,
VKBTG2009}) ). |
Flus-ratio anomalies observed. in multiple images of οσοquasars are often cited as evidence for substructures | Flux-ratio anomalies observed in multiple images of lensedquasars are often cited as evidence for substructures |
performances change including new physies. | performances change including new physics. |
In. brief. we run the same simulation switching cooling oll. | In brief, we run the same simulation switching cooling off. |
The result is shown in Fig 9. | The result is shown in Fig 9. |
Xs expected. without cooling. the eas component keeps round-shaped. whereas DM gets Hatter in the XZ plane (Carraro et al 1998a. Navarro White 1993). | As expected without cooling, the gas component keeps round-shaped, whereas DM gets flatter in the $X-Z$ plane (Carraro et al 1998a, Navarro White 1993). |
As for the adiabatie collapse we run a series of simulations of the formation of a disk galaxy at increasing number of We show the results in term of scalability ancl load balance in Figs 10. 11 and 12. | As for the adiabatic collapse we run a series of simulations of the formation of a disk galaxy at increasing number of We show the results in term of scalability and load balance in Figs 10, 11 and 12. |
In Fig. | In Fig. |
ll we present the scalability of 3 different code. sections. namely the gravity computation. the neighbors searching ancl the SPILL. measured. as the wall-clock time spent in a certain subroutine normalized to the global time-step. at increasing number of processors. | 11 we present the scalability of 3 different code sections, namely the gravity computation, the neighbors searching and the SPH, measured as the wall-clock time spent in a certain subroutine normalized to the global time-step, at increasing number of processors. |
Looking at this figure. we conclude hat the inclusion of the cooling processes does not. alfect he scalability. | Looking at this figure, we conclude that the inclusion of the cooling processes does not affect the scalability. |
Phe most significant deviation is visible in he gravity part of the code. like in the case of the adiabatic collapse (see Section 6). where the scaling starts cleviating significantly [rom the ideal one when using more than 32 On the other hand. the SPIE and neighbors searching xwis of the code scale. very well up to 64. processors. | The most significant deviation is visible in the gravity part of the code, like in the case of the adiabatic collapse (see Section 6), where the scaling starts deviating significantly from the ideal one when using more than 32 On the other hand, the SPH and neighbors searching parts of the code scale very well up to 64 processors. |
The parallel over-head. which measures the time spent to | The parallel over-head, which measures the time spent to |
ISO to be confusion limited at 90 pin and 175 pin by galaxies and galactic cirrus emission and hence this survey should be the deepest. PLR survey. possible with the satellite. | ISO to be confusion limited at 90 $\mu m$ and 175 $\mu m$ by galaxies and galactic cirrus emission and hence this survey should be the deepest FIR survey possible with the satellite. |
The area covered in the ELALS survey is ~13 square degrees at 15 ancl 00 microns. 7 square degrees at 6.7 microns and ~3 square degrees at. I75microns. | The area covered in the ELAIS survey is $\sim 13$ square degrees at 15 and 90 microns, $\sim 7$ square degrees at 6.7 microns and $\sim 3$ square degrees at 175microns. |
The ELAIS survey is ~50 times deeper at. 5-205900 than LRAS. | The ELAIS survey is $\sim$ 50 times deeper at $\mu m$ than IRAS. |
Thus our survey will allow us to explore LRAS-like populations to higher redshift ancl possibly unveil new classes of objects or unexpected phenomena. | Thus our survey will allow us to explore IRAS-like populations to higher redshift and possibly unveil new classes of objects or unexpected phenomena. |
We expect to detect thousands of galaxies. many of which will be at high redshifts ancl undergoing vigorous star formation. | We expect to detect thousands of galaxies, many of which will be at high redshifts and undergoing vigorous star formation. |
The expected large number of high-z LR galaxies should. provide vital information about the star formation rate out to z=1 and. possibly earlier. | The expected large number of high-z IR galaxies should provide vital information about the star formation rate out to z=1 and possibly earlier. |
The spatial resolution of ISO. will be insullicient to properly identify optically faint objects. | The spatial resolution of ISO will be insufficient to properly identify optically faint objects. |
At 15 microns. the survey resolution is 10 aresee ancl at 90 microns it will be about one arc minute. | At 15 microns, the survey resolution is $\sim$ 10 arcsec and at 90 microns it will be about one arc minute. |
Complementary radio data will play a crucial role in identifving many of the most. interesting objects. as they did in the carly days of X-ray astronomy (c.g. (νο N-1) and in more recent times for LUXS (e.g. HAS L10214|4714 (Rowan-Robinson οἱ αἱ. | Complementary radio data will play a crucial role in identifying many of the most interesting objects, as they did in the early days of X-ray astronomy (e.g. Cyg X-1) and in more recent times for IRAS (e.g. IRAS F10214+4714 (Rowan-Robinson et al. |
1991). | 1991). |
In this paper we report the description. of the radio observations obtained in the three ISO-IELATIS survey regions in the northern celestial hemisphere 11610]5430. N2.11636|4115 ancl N3_11429|3306). | In this paper we report the description of the radio observations obtained in the three ISO-ELAIS survey regions in the northern celestial hemisphere 1610+5430, 1636+4115 and 1429+3306). |
The observations are made with the Verv Large Array (VLA) radio telescope at ΜΑ] (200m) in the VLA C-configuration (maximum baseline 11km) with a resolution (LNLIND of ~15 aresec. | The observations are made with the Very Large Array (VLA) radio telescope at 1.4GHz (20cm) in the VLA C-configuration (maximum baseline 11km) with a resolution (FWHM) of $\sim15$ arcsec. |
The aim of these VLA observation was to obtain an uniform covering of the ELALS regions with a rms noise limit of —50 μὴν. | The aim of these VLA observation was to obtain an uniform covering of the ELAIS regions with a rms noise limit of $\sim$ 50 $\mu$ Jy. |
These VLA observations will »o essential in the optical identification phase of the ELALS sources and in assessing the reliability of the ELAIS source ists. | These VLA observations will be essential in the optical identification phase of the ELAIS sources and in assessing the reliability of the ELAIS source lists. |
Aloreover. with a radio survey it will be possible to investigate the raclioflarinfrared correlation in star forming galaxies to [Lux densities deeper than those reached by LAS. | Moreover, with a radio survey it will be possible to investigate the radio/far–infrared correlation in star forming galaxies to flux densities deeper than those reached by IRAS. |
llelou. Sofier Rowan-Robinson (1985) noted a strong correlation between radio and far infrared [lux for star forming galaxies. valid over a very wide range of infrarec luminosities. and this has been confirmed. in many other studies (e.g. Wunderlich. IxKIein Wielebinski 1987: Condon. Anderson Lelou 1991). | Helou, Sofier Rowan-Robinson (1985) noted a strong correlation between radio and far infrared flux for star forming galaxies, valid over a very wide range of infrared luminosities, and this has been confirmed in many other studies (e.g. Wunderlich, Klein Wielebinski 1987; Condon, Anderson Helou 1991). |
The radio emission is interpretec as à svnchrotron radiation from relativistic electron. which rave leaked out of supernova remnants. | The radio emission is interpreted as a synchrotron radiation from relativistic electron which have leaked out of supernova remnants. |
It is expected tha us correlation should. extend: below the LRAS flux leve Lgince the majority of the sub-mJy radio sources have been --dentified. with faint blue galaxies whith spectra similar to rose Of star forming objects (Benn et al. | It is expected that this correlation should extend below the IRAS flux level since the majority of the sub-mJy radio sources have been identified with faint blue galaxies whith spectra similar to those of star forming objects (Benn et al. |
1993). | 1993). |
In addition. combining deep radio and optical data =vith the [SO survey fuses will provide information on 1 trivaviate L-racio-optical luminosity [function ancl its evolution and the contribution of starburst galaxies to the sub-niJv. racio source Counts. | In addition, combining deep radio and optical data with the ISO survey fluxes will provide information on the trivariate IR-radio-optical luminosity function and its evolution and the contribution of starburst galaxies to the sub-mJy radio source counts. |
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