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. Such a line ofile morphology is synonymous with absorption/emission far above anddetached from the photosphere. and usually associated with chemical stratification of the corresponding element or density Kinks: here it stems solely from an caused by time-dependence effects.
Such a line profile morphology is synonymous with absorption/emission far above and from the photosphere, and usually associated with chemical stratification of the corresponding element or density kinks; here it stems solely from an caused by time-dependence effects.
The lline is entirely absent in the steady-state counterpart model three weeks after the explosion because. in the relevant region. helium atoms are once ionized with a probability of only one in 10310,
The line is entirely absent in the steady-state counterpart model three weeks after the explosion because, in the relevant region, helium atoms are once ionized with a probability of only one in $^{20-30}$ .
In Fig. Il.
In Fig. \ref{fig_hei_pip}, ,
we reproduce Fig.
we reproduce Fig.
1O— and illustrate. the absorption/emission sites for uat 48.7 days in the time-dependent version of Model A. Absorption stems from the regions in the direction of the stellar disk. far from the photosphere which has a velocity of ~4500 on that day.
\ref{fig_balmer_pip} and illustrate the absorption/emission sites for at 48.7 days in the time-dependent version of Model A. Absorption stems from the regions in the direction of the stellar disk, far from the photosphere which has a velocity of $\sim$ on that day.
For impact rays not intersecting the photodisk. only emission occurs. again atf large distances above the photosphere. and over a small volume.
For impact rays not intersecting the photodisk, only emission occurs, again at large distances above the photosphere, and over a small volume.
The emission flux is small. appears as a flat-topped profile. and is thus barely noticeable above the continuum.
The emission flux is small, appears as a flat-topped profile, and is thus barely noticeable above the continuum.
Of this lline in the synthetic spectrum. one can only observe the strongly blueshifted absorption component. the location of the flat-topped emission overlapping with the much stronger P line at this late epoch.
Of this line in the synthetic spectrum, one can only observe the strongly blueshifted absorption component, the location of the flat-topped emission overlapping with the much stronger $\gamma$ line at this late epoch.
This is illustrated differently in Fig. 12..
This is illustrated differently in Fig. \ref{fig_montage_hei},
where we show a montage of synthetic spectra covering the 1—1.2j/m region and the entire evolution computed for Model A. In the left panel. we Pal10w the full synthetic spectrum. ie.. including all species.
where we show a montage of synthetic spectra covering the $\mu$ m region and the entire evolution computed for Model A. In the left panel, we show the full synthetic spectrum, i.e., including all species.
In the right panel. we use the converged CMFGEN model at each time in the sequence and solve the formal solution of the radiative transfer problem by including bound-bound transitions of helium only.
In the right panel, we use the converged CMFGEN model at each time in the sequence and solve the formal solution of the radiative transfer problem by including bound-bound transitions of helium only.
Note how at early times (dark curves) the profiles in the two panels match. indicating that helium is the main feature contributing in emission and absorption at early times.
Note how at early times (dark curves) the profiles in the two panels match, indicating that helium is the main feature contributing in emission and absorption at early times.
After three weeks. the emission is dominated by P and the absorption gives rise to a dip. detached from the P~ absorption. and shifting further to the blue (i.e.. in an opposite sense to the photosphere which recedes to smaller velocities. as well as toallotherline-protile absorptions).
After three weeks, the emission is dominated by $\gamma$ and the absorption gives rise to a dip, detached from the $\gamma$ absorption, and shifting further to the blue (i.e., in an opposite sense to the photosphere which recedes to smaller velocities, as well as toallotherline-profile absorptions).
with the enhanced conductive heating.
with the enhanced conductive heating.
Alternatively. it could be interpreted that injection of larger energy in (he corona can heat more deeply to the higher density chromosphere by ihe downward thermal conduction.
Alternatively, it could be interpreted that injection of larger energy in the corona can heat more deeply to the higher density chromosphere by the downward thermal conduction.
A combination of (he above two relations of eqs.(23)) and (24)) roughly give on given 7 and fas. Which reminds us of the famous RTV scaling law (Rosner. 1978).. Trasc1400(piLgi)3 . for. closed magnetic.κ loops.. where pi(dvn $a.7) is loop pressure and Ly (em) is loop height.
A combination of the above two relations of \ref{eq:factmx}) ) and \ref{eq:fwptr}) ) roughly give on given $\tau$ and $f_{\rm max}$, which reminds us of the famous RTV scaling law \citep{rtv78}, $T_{\rm max} \simeq 1400 (p_{\rm l} L_{\rm h})^{1/3}$ , for closed magnetic loops, where $p_{\rm l}$ (dyn $^{-2}$ ) is loop pressure and $L_{\rm h}$ (cm) is loop height.
Taking rp; instead of Lj lor our model. we can actually derive a relation between μας aud (pira) [rom our results as showing a form analogous (o that of the original RTV laa.
Taking $r_{\rm Tmax}$ instead of $L_{\rm h}$ for our model, we can actually derive a relation between $T_{\rm max}$ and $(p_{\rm tr} r_{\rm Tmax})$ from our results as showing a form analogous to that of the original RTV law.
This is because we consider the same energv balance among thermal conduction. radiative cooling. and heating with the same boundary condition that the conductive [αν should become almost zero at the base (eq.(17))). although the configurations are quite different (closed loop for the RTV law ancl open flow tube for ours).
This is because we consider the same energy balance among thermal conduction, radiative cooling, and heating with the same boundary condition that the conductive flux should become almost zero at the base \ref{eq:bc3}) )), although the configurations are quite different (closed loop for the RTV law and open flow tube for ours).
The slight discrepaucies of the prefactor and power-law index are caused by the [act that the REV. law was derived on the assumption of spatially uniform heating along the loop. while our heating function is determined by eq.(13)). which is nol uniform at all.
The slight discrepancies of the prefactor and power-law index are caused by the fact that the RTV law was derived on the assumption of spatially uniform heating along the loop, while our heating function is determined by \ref{eq:wvgr}) ), which is not uniform at all.
In fig.5 we show anticipated proton flux. (npo)iyc. at LAU as a function of Fy Tor different 7 and μι. with observational constraints. (πρ0εν=(8.841.5)x107m 25 1 (Shaded). compiled by Withbroe(1988) as the empirical value for «quet corona’ which is supposed to corresponds to the mid- to low-latitude region generating the low-speed wind.
In \ref{fig:fwmd} we show anticipated proton flux, $(n_{\rm p}v)_{\rm 1AU}$, at 1AU as a function of $F_{\rm w,0}$ for different $\tau$ and $f_{\rm max}$, with observational constraints, $(n_{\rm p}v)_{\rm 1AU}=(3.8\pm 1.5)\times 10^8$ $^{-2}$ $^{-1}$ (Shaded), compiled by \citet{wtb88} as the empirical value for 'quiet corona' which is supposed to corresponds to the mid- to low-latitude region generating the low-speed wind.
Larger 7 waves give a greater (Προ: owing to the effective transport of dissipated energy to solar wind flow by avoiding radiative escape.
Larger $\tau$ waves give a greater $(n_{\rm p}v)_{\rm 1AU}$ owing to the effective transport of dissipated energy to solar wind flow by avoiding radiative escape.
Introduction of the non-radial expansion of the flow tube reduces (Προ)εν. because input energy per unit flow (tube normalized ab LAU decreases on increasing μις. even though one inputs identical wave energy flux al (he inner corona.
Introduction of the non-radial expansion of the flow tube reduces $(n_{\rm p}v)_{\rm 1AU}$, because input energy per unit flow tube normalized at 1AU decreases on increasing $f_{\rm max}$, even though one inputs identical wave energy flux at the inner corona.
Therefore. (he larger areal expansion straight(lorwarclly reduces mass flux of the solar wind.
Therefore, the larger areal expansion straightforwardly reduces mass flux of the solar wind.
As to dependences on Fig. (Προ). has an upper limit for given T and fray. (
As to dependences on $F_{\rm w,0}$, $(n_{\rm p}v)_{\rm 1AU}$ has an upper limit for given $\tau$ and $f_{\rm max}$. (
νου models of 7= 300sand fia,=1 are supposed to have the upper limit ina Fyy>2x I0?erg em7s ! area.)
Even models of $\tau=300$ sand $f_{\rm max}=1$ are supposed to have the upper limit in a $F_{\rm w,0}>2\times 10^6$ erg $^{-2}$ $^{-1}$ area.)
Figure 5. shows that only one case emploving G.fuas)=(800.1) can reproduce observed. proton flux of the slow wind.
Figure \ref{fig:fwmd} shows that only one case employing $(\tau,f_{\rm max}) =(300,1)$ can reproduce observed proton flux of the slow wind.
Unfortunately
Unfortunately
instantaneously Since we are interested in the effects on dust enrichment. we compare our results with the dust abundance in galaxies.
instantaneously Since we are interested in the effects on dust enrichment, we compare our results with the dust abundance in galaxies.
Given that the dust enrichment is closely related with the metal enrichment. it is convenient to derive the relation between dust-to-gas ratio and metallicity (e.g.Lisenfeld&Ferrara1998).
Given that the dust enrichment is closely related with the metal enrichment, it is convenient to derive the relation between dust-to-gas ratio and metallicity \citep[e.g.][]{lisenfeld98}.
. The dust mass is usually estimated by the far-infrared emission to trace the emission from large grains. which occupy a significant fraction of the dust mass.
The dust mass is usually estimated by the far-infrared emission to trace the emission from large grains, which occupy a significant fraction of the dust mass.
The dust mass estimated from the far-infrared emission can miss very small grains and very cold dust.
The dust mass estimated from the far-infrared emission can miss very small grains and very cold dust.
The contribution from very small grains to the dust mass is not significant2011).
The contribution from very small grains to the dust mass is not significant.
Very cold dust traced in longer wavelengths than submillimetre may have a significant contribution to the total dus mass especially in metal-poor galaxies (Gallianoetal.2003.2005:Galametzetal.201 L3... but its abundance is significantly affected by the assumed emissivity index of large grains.
Very cold dust traced in longer wavelengths than submillimetre may have a significant contribution to the total dust mass especially in metal-poor galaxies \citep{galliano03,galliano05,galametz11}, but its abundance is significantly affected by the assumed emissivity index of large grains.
If we miss the contribution from very cold dust. the observational dust-to-gas ratio is underestimated in this paper. which enhances the importance of the dust growth in clouds to explain the additional contribution from very cold dust.
If we miss the contribution from very cold dust, the observational dust-to-gas ratio is underestimated in this paper, which enhances the importance of the dust growth in clouds to explain the additional contribution from very cold dust.
Although there is a large variety in the grain size distribution derived from the dust emission spectrum (Gallianoetal.2005). the regulation mechanism of grain size distribution in the ISM is not fully understood.
Although there is a large variety in the grain size distribution derived from the dust emission spectrum \citep{galliano05}, the regulation mechanism of grain size distribution in the ISM is not fully understood.
Therefore. in this paper. we examine various grain size distributions 6 function with various typical sizes and power law with various power indeces).
Therefore, in this paper, we examine various grain size distributions $\delta$ function with various typical sizes and power law with various power indeces).
Sputtering and shattering in shocks (Jonesetal.1996). and shattering in interstellar turbulenceSN (Yan.Lazarian.&Draine2004:Hirashitaet may play a signiticant role in determining the shape of grain size distribution.
Sputtering and shattering in SN shocks \citep{jones96} and shattering in interstellar turbulence \citep*{yan04,hirashita10} may play a significant role in determining the shape of grain size distribution.
Inclusion of these processes into the evolution of grain size distribution is left for future work to concentrate on the grain growth in clouds in this paper.
Inclusion of these processes into the evolution of grain size distribution is left for future work to concentrate on the grain growth in clouds in this paper.
A simultaneous treatment of dust formation and destruction can be seen in Yamasawaetal.OLI)..
A simultaneous treatment of dust formation and destruction can be seen in \citet{yamasawa11}.
Photo-processes are neglected in this paper.
Photo-processes are neglected in this paper.
Dust may contain ice mantle when it is injected into the ISM from clouds.
Dust may contain ice mantle when it is injected into the ISM from clouds.
When the dust is exposed to stellar radiation. the evaporation of the ice mantle may lead to grain disaggregation.
When the dust is exposed to stellar radiation, the evaporation of the ice mantle may lead to grain disaggregation.
If small refractory grains are ejected in the disaggregation. the total amount of refractory dust (silicate and graphite in this paper) does not change. so that our conclusion below is not affected.
If small refractory grains are ejected in the disaggregation, the total amount of refractory dust (silicate and graphite in this paper) does not change, so that our conclusion below is not affected.
Tf the dust is destroyed and the refractory elements are returned into the gas phase in the photo-processing. the dust abundance may be affected by this process.
If the dust is destroyed and the refractory elements are returned into the gas phase in the photo-processing, the dust abundance may be affected by this process.
The destruction of dust by photo-processing. if any. ean be effectively included into the dust destruction efficiency (Oy defined later). but we will see later that only the destruction by SNe is enough to explain the dust-to-metal ratio (Section 5.2)).
The destruction of dust by photo-processing, if any, can be effectively included into the dust destruction efficiency $\beta_\mathrm{SN}$ defined later), but we will see later that only the destruction by SNe is enough to explain the dust-to-metal ratio (Section \ref{subsec:Zcr}) ).
We detine the increased fraction of dust mass in a cloud. 7. which is evaluated as where 7.) is the lifetime of clouds hosting the grain growth and equation (18)) is used from the second to the third step.
We define the increased fraction of dust mass in a cloud, $\beta$, which is evaluated as where $\tau_\mathrm{cl}$ is the lifetime of clouds hosting the grain growth and equation \ref{eq:a3_expansion}) ) is used from the second to the third step.
The dust mass in the cloud becomes (.7|T) times the initial value after the cloud lifetime. when the dust grown in the cloud returns in the diffuse ISM (see equation 189).
The dust mass in the cloud becomes $(\beta +1)$ times the initial value after the cloud lifetime, when the dust grown in the cloud returns in the diffuse ISM (see equation \ref{eq:a3_expansion}) ).
Now we consider how to implement our results into dust enrichment models for an entire galactic system.
Now we consider how to implement our results into dust enrichment models for an entire galactic system.
By using 7. the increasing rate of dust mass by accretion in clouds should. be written as where Afi: is the total dust mass in the galaxy. and Vo is the mass fraction of clouds hosting the grain growth to the total gas mass.
By using $\beta$, the increasing rate of dust mass by accretion in clouds should be written as where $M_\mathrm{dust}$ is the total dust mass in the galaxy, and $X_\mathrm{cl}$ is the mass fraction of clouds hosting the grain growth to the total gas mass.
In equation (29)). we have assumed that the time-scale of dust enrichment is much longer than the cloud lifetime so that the mass increasing rate in each cloud is estimated by «7/71 times the dust mass in the cloud.
In equation \ref{eq:dMdt2}) ), we have assumed that the time-scale of dust enrichment is much longer than the cloud lifetime so that the mass increasing rate in each cloud is estimated by $\beta /\tau_\mathrm{cl}$ times the dust mass in the cloud.
On the other hand. Hirashita(20002) express the increasing rate of the dust mass in a galaxy through the grain growth in elouds by where Tara. 18 the growth time-scale of the dust in a cloud. and € is the fraction of metals in gas phase (Section 22)).
On the other hand, \citet{hirashita00} express the increasing rate of the dust mass in a galaxy through the grain growth in clouds by where $\tau_\mathrm{grow}$ is the growth time-scale of the dust in a cloud, and $\xi$ is the fraction of metals in gas phase (Section \ref{subsec:growth}) ).
A similar expression for the increasing rate of the dust mass is widely adopted in dust enrichment models 2008).
A similar expression for the increasing rate of the dust mass is widely adopted in dust enrichment models \citep[e.g.][]{dwek98,inoue03,calura08}. .
At rycor there should be no energv loss because the initial energy is deposited in a volume of linear size r.
At $r_B\sim r$ there should be no energy loss because the initial energy is deposited in a volume of linear size $r$.
Eq.
Eq.
5. reflects this.
\ref{destruct_cond_shock} reflects this.
Then the ri, expression in Eq.
Then the $r_{\rm eq}$ expression in Eq.
3. is still valid.
\ref{destruct_cond} is still valid.
The above calculation of the size spectrin (reals V(r) as constant in Gime.
The above calculation of the size spectrum treats $N(r)$ as constant in time.
To maintain is steady-state exactly would require the power law to extend to bodies of infinite size. Evhich is impossible.
To maintain this steady-state exactly would require the power law to extend to bodies of infinite size, which is impossible.
To find (he range of masses where (his assumption holds. we first find je SIZE rig ol the largest ΙΟ (to have experienced a destructive collision alter an elapsed ime r.
To find the range of masses where this assumption holds, we first find the size $r_{\rm break}$ of the largest KBO to have experienced a destructive collision after an elapsed time $\tau$.
We equate (he (timescale for destructive collisions for each NBO of size ri (o 7 —ising Eqs.
We equate the timescale for destructive collisions for each KBO of size $r_{\rm break}$ to $\tau$ using Eqs.
1 and 5..
\ref{steadystate} and \ref{destruct_cond_shock}.
To get Vr) we note that bodies of size r>rijs. having never collided. V.Yould be effectivelv primordial at time 7.
To get $N(r)$ we note that bodies of size $r>r_{\rm break}$, having never collided, should be effectively primordial at time $\tau$.
For their size spectrum we write V(r)=Nor!© where Vyc4x107em1 from observations (2)..
For their size spectrum we write $N(r)=N_0 r^{1-q_0}$ where $N_0\sim 4\times 10^{7q_0-3}\;\cm ^{q_0-1}$ from observations \citep{trujillo01}.
This is equivalent to a Kuiper belt with 4x10! bodies larger than 100 km.
This is equivalent to a Kuiper belt with $4\times 10^4$ bodies larger than 100 km.
Thev are spread over an area A1200AU? in the plane of the solar system (?).. so V.cο where ο=0.022vr| is the typical orbital angular velocity of the Ixuiper belt.
They are spread over an area $A\simeq 1200\;{\rm AU}^2$ in the plane of the solar system \citep{trujillo01}, so $V\simeq Av/\Omega$ where $\Omega=0.022\;{\rm yr}^{-1}$ is the typical orbital angular velocity of the Kuiper belt.
With 4 lor the slope below the break and. as above. qu and Ny for the slope and normalization above the break. we have If we set724.5x10?vr to be the agec» of the solar svstem. take 3/2/<9«3. and use the observed gy=5. we get This is consistent with the observed break position of ~70 km (?)..
With $q$ for the slope below the break and, as above, $q_0$ and $N_0$ for the slope and normalization above the break, we have If we set$\tau\simeq 4.5\times 10^9\;{\rm yr}$ to be the age of the solar system, take $3/2<\beta<3$, and use the observed $q_0=5$, we get This is consistent with the observed break position of $\sim$ 70 km \citep{bernstein03}.
Note that if the svstem had had (he high velocity dispersion assumed above over a time considerably shorter than 4.5 Gvr. the break would have occurred al a much smaller ABO size.
Note that if the system had had the high velocity dispersion assumed above over a time considerably shorter than 4.5 Gyr, the break would have occurred at a much smaller KBO size.
We therefore inler (hat the Kuiper belts current excited state has been a lone-lived phase of al least a few billion νους duration rather than a recent phenomenon.
We therefore infer that the Kuiper belt's current excited state has been a long-lived phase of at least a few billion years' duration rather than a recent phenomenon.
The evolution of the total mass and velocity dispersion of the Kuiper belt is a potential concern. as (he break location depends strongly on both.
The evolution of the total mass and velocity dispersion of the Kuiper belt is a potential concern, as the break location depends strongly on both.
The mass of the Ixuiper belt may have been larger by a [actor of ~L00 when the solar svstem was very voung (10*—LO vears old) (see.lovexample.?)..
The mass of the Kuiper belt may have been larger by a factor of $\sim$ 100 when the solar system was very young $10^7-10^8$ years old) \citep[see, for example,][]{kenyon02}.
The collision frequency would have been much higher (hen. so collisions during that period might be expected to have increased the break radius.
The collision frequency would have been much higher then, so collisions during that period might be expected to have increased the break radius.
At that lime. though. the velocity dispersion of INBO precursors is believed to have been just 1 m/s (see.forexample. 2)..
At that time, though, the velocity dispersion of KBO precursors is believed to have been just $\sim$ 1 m/s \citep[see, for example,][]{goldreich02}. .
With this impact velocity. r471 km.so only targets of size «1 ki
With this impact velocity, $r_{\rm eq} \sim 1\;\km$ ,so only targets of size $<$ 1 km
With this impact velocity. r471 km.so only targets of size «1 kin
With this impact velocity, $r_{\rm eq} \sim 1\;\km$ ,so only targets of size $<$ 1 km
The lens is made of a point mass. a constant shear. aid a whole variety of multipoles.
The lens is made of a point mass, a constant shear, and a whole variety of multipoles.
From (he index of the vector field wilh zeros ancl poles. it is obtained (hal in—5».=—2where n. is the number of positive/negative images.
From the index of the vector field with zeros and poles, it is obtained that $n_+-n_-=-2$where $n_\pm$ is the number of positive/negative images.
Thus the number of images is even.
Thus the number of images is even.
For w=x. there are four images. one al 2=x and three degenerateimages al 2=0.
For $\omega = \infty$, there are four images, one at $z=\infty$ and three degenerateimages at $z=0$.
As w moves toward the lenses. the three degenerate images individualizes.
As $\omega$ moves toward the lenses, the three degenerate images individualizes.
It is expected that the number of images is four outside the caustic and six inside so that n.--η=—2.
It is expected that the number of images is four outside the caustic and six inside so that $n_+-n_-=-2$.
Itis known that the original lens equation (eq.(2)) or(2))) without approximations produces four or six images (ErcllandSchneider1993:Petters1997).. and RBIO succeeded in deriving (he sixth order analvtic polvnonual equation [rom the lens equaiton.
Itis known that the original lens equation \ref{eqLeqOne}) ) or\ref{eqLeqTwo}) )) without approximations produces four or six images \citep{ES93, petters}, and RB10 succeeded in deriving the sixth order analytic polynomial equation from the lens equaiton.
However. the approximate lens equation (3.2.1)) has a third order pole. and we will see (hat (wo images are dark images remaining near the pole.
However, the approximate lens equation \ref{eqApproxLeqTwo}) ) has a third order pole, and we will see that two images are dark images remaining near the pole.
We reler to them as ignorable images.
We refer to them as ignorable images.
Jacobian matrix components are : [=O0e=1+fe: g = = 3-7 — +
Jacobian matrix components are f = = 1 +; g = = - - + +.
esL..(44) Set 2=re" and the Jacobian determinant J=||?>—3 can be computed up to the second order.
Set $z = r e^{i\theta}$ and the Jacobian determinant $J = |f|^2 -|g|^2$ can be computed up to the second order.
In the linear order. Jiu =1— ((1 — 4")). and the critical curve. ρω= 0. is given by 8— zu)
J = 1 - + + - ) - - - In the linear order, = 1 - (1 - ), and the critical curve, $J_{linear}=0$ , is given by = r - r^5 ; r 0 .
The discovery that approximately half the energy ever radiated by galaxies is received on Earth in the far-infrared waveband (Puget et 11996; Fixsen et 11998) implies. that galaxies must show strong evolution that is hidden from optical telescopes (Gispert. Lagache Puget 2000).
The discovery that approximately half the energy ever radiated by galaxies is received on Earth in the far-infrared waveband (Puget et 1996; Fixsen et 1998) implies that galaxies must show strong evolution that is hidden from optical telescopes (Gispert, Lagache Puget 2000).
A decade ago. deep surveys with the SCUBA camera on the James Clerk Maxwell Telescope resolved much of the far-infrared background (FIRB) at 850 um into individual sources (Barger et 11998: Hughes et 11998).
A decade ago, deep surveys with the SCUBA camera on the James Clerk Maxwell Telescope resolved much of the far-infrared background (FIRB) at 850 $\mu$ m into individual sources (Barger et 1998; Hughes et 1998).
These sources are mostly extremely luminous dust-enshrouded galaxies at z>2 (Chapman et 22005) with an average implied star-formation rate (1f the ultimate source of the energy is star formation) of ~400 Μ..year”! (Coppin et 22006). much greater than the star-formation rates m galaxies like our own.
These sources are mostly extremely luminous dust-enshrouded galaxies at $\rm z > 2$ (Chapman et 2005) with an average implied star-formation rate (if the ultimate source of the energy is star formation) of $\simeq$ 400 $\rm M_{\odot}\ year^{-1}$ (Coppin et 2006), much greater than the star-formation rates in galaxies like our own.
However. the energy density in the FIRB at 850 jim is =30 times less than at 200 uim where the FIRB is at à maximum. and both the spectral shape of the FIRB and statistical ‘stacking’ analyses (Dole et al.
However, the energy density in the FIRB at 850 $\mu$ m is $\simeq$ 30 times less than at 200 $\mu$ m where the FIRB is at a maximum, and both the spectral shape of the FIRB and statistical `stacking' analyses (Dole et al.
2006: Pascale et 22009) imply that much of the FIRB is actually produced by sources at lower redshift (Gispert et 22000: Dole et al.
2006; Pascale et 2009) imply that much of the FIRB is actually produced by sources at lower redshift (Gispert et 2000; Dole et al.
2006: Pascale et 22009).
2006; Pascale et 2009).
The launch of the Herschel Space Observatory (Pilbratt et 22010) in May 2009 has given us the opportunity to resolve a significant fraction of the FIRB at wavelengths where its energy density is at a maximum.
The launch of the Herschel Space Observatory (Pilbratt et 2010) in May 2009 has given us the opportunity to resolve a significant fraction of the FIRB at wavelengths where its energy density is at a maximum.
In this letter. using the first data from the Herschel Multi-tiered Extragalactic Survey (HerMES: Oliver et al..
In this letter, using the first data from the Herschel Multi-tiered Extragalactic Survey (HerMES; Oliver et al.,
in prep). we investigate the evolution implied by the existence of the FIRB by measuring the evolution of the galaxy luminosity function at 250 µη. We everywhere assume a standard concordance cosmology: Qn=0.28.Q420.72.Ho72kms!Μρο.
in prep), we investigate the evolution implied by the existence of the FIRB by measuring the evolution of the galaxy luminosity function at 250 $\mu$ m. We everywhere assume a standard concordance cosmology: $\rm \Omega_M=0.28,\ \Omega_{\Lambda}=0.72,\ H_0 = 72\ km\ s^{-1}\ Mpc^{-1}$.
The images that we analyse in this letter were taken at 250 jm with the SPIRE instrument onHerschel. whose m-orbit performance and scientific capabilities are described in Griffin et (2010).
The images that we analyse in this letter were taken at 250 $\mu$ m with the SPIRE instrument on, whose in-orbit performance and scientific capabilities are described in Griffin et (2010).
The calibration methods and accuracy of SPIRE are described by Swinyward et ((2010).
The calibration methods and accuracy of SPIRE are described by Swinyward et (2010).
The three images consist of a shallow image of the Lockman Hole (LH) and deep images of the northern field of the Great Observatories Origins Deep Survey (GOODS-North) and of a field within the Lockman Hole (LH-North).
The three images consist of a shallow image of the Lockman Hole (LH) and deep images of the northern field of the Great Observatories Origins Deep Survey (GOODS-North) and of a field within the Lockman Hole (LH-North).
For the latter two images the dominant source of noise is confusion due to numerous faint sources. which is 5.8 mJy beam' at 250 jm (Nguyen et 22010).
For the latter two images the dominant source of noise is confusion due to numerous faint sources, which is 5.8 mJy $^{-1}$ at 250 $\mu$ m (Nguyen et 2010).