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The optical colors of the galaxy. (corrected for the Milky Wav absorption aud. A-correction) are rather blue (5| 2u055.VR= 0.55) and are typical for a late-type or ireeular galaxy. | The optical colors of the galaxy (corrected for the Milky Way absorption and $K$ -correction) are rather blue $B-V=0.55$ , $V-R=0.35$ ) and are typical for a late-type or irregular galaxy. |
A small edge-on (av be background) ealaxy with B=19.5. BYo=09. VvRo=10.6 (observational values) is located 19" SW of the PCC 100092 uucleus. | A small edge-on (may be background) galaxy with $B=19.5$, $B-V=0.9$, $V-R=+0.6$ (observational values) is located $''$ SW of the PGC 400092 nucleus. |
The southern galaxy of the triplet (PGC 399718) possess a bar and a two-arlus spiral structure. | The southern galaxy of the triplet (PGC 399718) possess a bar and a two-arms spiral structure. |
The BoV. color of the galaxy looks unusually red for spirals aud his can be an indication of a high inclination. | The $B-V$ color of the galaxy looks unusually red for spirals and this can be an indication of a high inclination. |
A hieh velocity exadieut at P.A.=90° (see sect. | A high velocity gradient at $^{\rm o}$ (see sect. |
3.2) also supports lis poiut of view. | 3.2) also supports this point of view. |
Therefore. the galaxy structure may ος not planar - with almost edec-ou inner part aud more ACce-on outer ones. | Therefore, the galaxy structure may be not planar - with almost edge-on inner part and more face-on outer ones. |
The simulations were performed with a N-body particle code which inclides eas clvnamics and star formation (see BCO3). | The simulations were performed with a $N$ -body particle code which includes gas dynamics and star formation (see BC03). |
The main characteristics of this code is that the potential is computed by FFT ou a 3D Cartesian eric. | The main characteristics of this code is that the potential is computed by FFT on a 3D Cartesian grid. |
The size of the evid is 1287. and the Fourier images are suppressed by the method of James (1977). | The size of the grid is $^3$, and the Fourier images are suppressed by the method of James (1977). |
The softening leugth is equal to the size of a ericd cell. which is of the order of 1 to 3 kpc. according to the models. | The softening length is equal to the size of a grid cell, which is of the order of 1 to 3 kpc, according to the models. |
The particles. stars. gas and dark matter. are in total umber between 2 and 6.100, | The particles, stars, gas and dark matter, are in total number between 2 and $\cdot$ $^5$. |
The eas dvuamiics is modeled by the sticky-particles schiena. | The gas dynamics is modeled by the sticky-particles scheme. |
There are too many free pariueters dun a nunierica model to fit the morphology of AM. 1931-563 in details (sce discussion in DCO3). | There are too many free parameters in a numerical model to fit the morphology of AM 1934-563 in details (see discussion in BC03). |
. Therefore. we decided to fiuc a solution which describes reasonably only the genera observational properties of the galaxy. | Therefore, we decided to find a solution which describes reasonably only the general observational properties of the galaxy. |
We have considere wo scenarios for the polar ring formation: 1) ninor merecr with a simall companion. where the ring is formed by he tidal disruption of the simallest ealaxw and 2) fida accretion of matter from a imiassive. gas rich donor galaxy o the polar ring (as in DC). | We have considered two scenarios for the polar ring formation: 1) minor merger with a small companion, where the ring is formed by the tidal disruption of the smallest galaxy and 2) tidal accretion of matter from a massive, gas rich donor galaxy to the polar ring (as in BC03). |
The major merecr scenario (sce Introduction) has not been simulated. but models of lis scenario cau be found in Bekki (1998) and BCOS. | The major merger scenario (see Introduction) has not been simulated, but models of this scenario can be found in Bekki (1998) and BC03. |
Ta the framework of the first. scenario. we have cousicdered a minor imerecr with an iuclined orbit. where the rine results from the disuptiou of the comipauion. | In the framework of the first scenario, we have considered a minor merger with an inclined orbit, where the ring results from the disruption of the companion. |
Iu the preseut case. the mass of the companion should be of the mass of the lost galaxy: more uassive conripanions would disturb the host ealaxv too uuch and fori a SO-like object (Bournaud et al. | In the present case, the mass of the companion should be of the mass of the host galaxy: more massive companions would disturb the host galaxy too much and form a S0-like object (Bournaud et al. |
2001. 2005a). less massive ones can only form rines tha are nuch fainter than in the observed system. | 2004, 2005a), less massive ones can only form rings that are much fainter than in the observed system. |
Ou le basis of our observations. the main galaxy stellar mass was set Qo STUES=210H M. | On the basis of our observations, the main galaxy stellar mass was set to $_{stars}=2 \times 10^{11}$ $_{\sun}$. |
The gas fraction was fixed as Moja Maga 0.2. | The gas fraction was fixed as $_{gas}$ $_{stars}$ =0.2. |
We asstuned a spherical dark latter ido. described by a Plinuner sphere of scale-leugtl 30 μεoc and truncation radius 60 kpc. with a dark-to-visible nass ratio Mioyy Mas 70.6 msde the stellar radius. | We assumed a spherical dark matter halo, described by a Plummer sphere of scale-length 30 kpc and truncation radius 60 kpc, with a dark-to-visible mass ratio $_{DH}$ $_{stars}$ =0.6 inside the stellar radius. |
This nass was chosen to fit the rotation curve amplitude. | This mass was chosen to fit the rotation curve amplitude. |
The characteristic size of the exponential disk is taken from. he surface brightucss fit. | The characteristic size of the exponential disk is taken from the surface brightness fit. |
We have simulated several imuinor mergers with a COMpaluou nass of the primary nass. | We have simulated several minor mergers with a companion mass of the primary mass. |
Such a mass ratio is required to forma a massive chough rine. without disturbing the host disk more than iuthe observed. case. | Such a mass ratio is required to form a massive enough ring, without disturbing the host disk more than inthe observed case. |
The internal characteristics of the companion are fixed as Mo4;/ gu 70.25. aud Mp | Αςgas=. aud orbital xuaneters even in Table 5. | The internal characteristics of the companion are fixed as $_{gas}$ $_{stars}$ =0.25 and $_{DH}$ $_{stars+gas}$ =1.5, and orbital parameters given in Table 5. |
The result of two | The result of two |
the mean angular displacement oBIET mrces from NGC 1068 (Table 2. col 6) to calc"late the proper motion. | the mean angular displacement of its sources from NGC 1068 (Table 2, col 6) to calculate the proper motion. |
This is found to be <"n nailliarcsec per vr which would be difficult o detect even if observations were spaced. over niulv ye"WAS. | This is found to be $<0.7$ milliarcsec per yr which would be difficult to detect even if observations were spaced over many years. |
ο) Progenitor Configuration. | g) Progenitor Configuration. |
This paper has eximined ti6 position aud redshift data of the conrpact objects rear NGC LOGS as they curentv exist and no attempt has been iade to determine their configuration prior to ejecion. | This paper has examined the position and redshift data of the compact objects near NGC 1068 as they currently exist and no attempt has been made to determine their configuration prior to ejection. |
Alhoug jidtfo may be prenature to specula con what the progenitor source coufiguratiou πιο, iwe been. twre are some thines that aro worth notiic. | Although it may be premature to speculate on what the progenitor source configuration might have been, there are some things that are worth noting. |
In Fie. | In Fig. |
dtfi6 posiion angle of the dred sources (e) cun 0 seen to rotate sanoothly from A to D. ending up approximately )rpendieulu το the preseut position of the rotation axis and aligned with hne major axis of hne ceural torus. | 4 the position angle of the paired sources $\omega)$ can be seen to rotate smoothly from A to D, ending up approximately perpendicular to the present position of the rotation axis and aligned with the major axis of the central torus. |
If this change is produced by a sumnultaucous chauge iu he tip of the rotation axis with time the present posiion angeles of the udrs are at least consistewt with the possibility hat they all originally orbited im the major axis aue of the galaxy. | If this change is produced by a simultaneous change in the tip of the rotation axis with time the present position angles of the pairs are at least consistent with the possibility that they all originally orbited in the major axis plane of the galaxy. |
Since the paired sources ave equal-and-opposit¢ spectral shifts (z, 4). it is argued here that prior o ejection cach pair was a Lorbiting binary. | Since the paired sources have equal-and-opposite spectral shifts $_{\rm r-b}$ ), it is argued here that prior to ejection each pair was an orbiting binary. |
Note hat the z,ji, values (as high as z = 0.37 in pair Dj are too large to be exdlained as components of the modest ejecjon velocities aud seeni more ikcly to ο residual orbial velocities. | Note that the $_{\rm r-b}$ values (as high as z = 0.37 in pair D) are too large to be explained as components of the modest ejection velocities and seem more likely to be residual orbital velocities. |
Also. in Fie. | Also, in Fig. |
| the etters 73 or "b uext to the paire sources Indicates wjether he source is redshiftcc or blueshifted relative to the mean pair redshift. | 4 the letters "r" or "b" next to the paired sources indicates whether the source is redshifted or blueshifted relative to the mean pair redshift. |
The blueshifted pair members (1.2.L5) all lie wes of the redshitted nembers (3.5GALL). and any proposed xoesenitor configuration mst be able to explain lus. | The blueshifted pair members (1,2,4,8) all lie west of the redshifted members (3,5,6,11), and any proposed progenitor configuration must be able to explain this. |
It is of interes to note that if the displacements of 10 padr niklpoiits from the ealaxy are removed. al the bluesufted objects fal ou the West of the eaaxv. and the redshifted oues ou the East. | It is of interest to note that if the displacements of the pair midpoints from the galaxy are removed, all the blueshifted objects fall on the West of the galaxy, and the redshifted ones on the East. |
This seeus to inply tha if the pairs were initially iu orbit. hey all orjted 1u the same cirection. | This seems to imply that if the pairs were initially in orbit, they all orbited in the same direction. |
h) Hidden bhieshitts. | h) Hidden blueshifts. |
It is worth pointing out that t1e alo"e ejectiou model demonstrates how large bheshifted compoucuts [neeative zo on) values] can be present but still be conipletelv camouflaged bv even larger iutriusic redshifts. | It is worth pointing out that the above ejection model demonstrates how large blueshifted components [negative $_{\rm r-b}$ ) values] can be present but still be completely camouflaged by even larger intrinsic redshifts. |
1) Redshifts of sources 10 aud 12. | i) Redshifts of sources 10 and 12. |
Until the recshifts of sources LO aud 12 can be lcasured it wi] not be possible to determine the size of the iutrisic Component of zia suce this redshift compoient will also contain a component due to the N-S. -o-s triplet ejection velocity (if 5+n). | Until the redshifts of sources 10 and 12 can be measured it will not be possible to determine the size of the intrinsic component of $_{\rm mean}$, since this redshift component will also contain a component due to the N-S, l-o-s triplet ejection velocity (if $\gamma \neq 0$ ). |
However. tlat component can be estimated Oo be at least as high as z — 0.6 (for >20). | However, that component can be estimated to be at least as high as z = 0.6 (for $\gamma > 0$ ). |
These results «o allow the redshift of source 12 o be roughly estimated for a future test of the uodel. | These results do allow the redshift of source 12 to be roughly estimated for a future test of the model. |
Since xmrcees LO aud 12 are both moving oward us iut 1ο model. their redshifts must be ess than the μι values of their respective pairs. | Since sources 10 and 12 are both moving toward us in the model, their redshifts must be less than the $_{\rm mean}$ values of their respective pairs. |
If the intrinsic redshift component of zii ds equal to that oits associated sineet. as assed above, then the redshifts of 10 ux 12 should differ from z= 1.2 and 1.1. respecively. bv. twice heir Lo-s ejection velocities. | If the intrinsic redshift component of $_{\rm mean}$ is equal to that of its associated singlet, as assumed above, then the redshifts of 10 and 12 should differ from z = 1.2 and 1.1, respectively, by twice their l-o-s ejection velocities. |
This difference will ve sinallif the radial ejection velocities are similar o those found for the sinelets iun triets A aud D. The vedshitt of source 12 is estimated here to veg = O.6O340.1 using the estimates triplet age and the angular separation of the sinelet from. he rotation axis. | This difference will be small if the radial ejection velocities are similar to those found for the singlets in triplets A and B. The redshift of source 12 is estimated here to be z = $\pm0.1$ using the estimated triplet age and the angular separation of the singlet from the rotation axis. |
No atteupt has been made to estimate a value for the redshift of source 10 j)ecause the aneular distances involved are too sunall to estimate accurately. | No attempt has been made to estimate a value for the redshift of source 10 because the angular distances involved are too small to estimate accurately. |
Tt the four triplets defined by the boxes in Fig. | If the four triplets defined by the boxes in Fig. |
l. whose component sources a] he at a simular distance foun NGC 1068. have been ejected from the ealaxy in separate events. it has Όσοι s1OW]. that the preseut positious of f1C πο.ο ο1 the sky can lead to a detailed. reasonable. iterally consistent ejection model. | 1, whose component sources all lie at a similar distance from NGC 1068, have been ejected from the galaxy in separate events, it has been shown that the present positions of the sources on the sky can lead to a detailed, reasonable, internally consistent ejection model. |
For tje source position and redshift data to lead to the sale sinele-pair μμ...ure in cach of the four trijets. that roates sx1))0thlv about the rotation axis in the same direcion that the ealaxy rotates. aud which is. at the same time. correlated with the triplet size and age. is remarkable. | For the source position and redshift data to lead to the same singlet-pair structure in each of the four triplets, that rotates smoothly about the rotation axis in the same direction that the galaxy rotates, and which is, at the same time, correlated with the triplet size and age, is remarkable. |
For the triplet ejection direcjon to lie along the rotation axis of the ealaxy. the one direction most closely associated with he ejection of matter from and active galaxy. In οςnally remarkable. | For the triplet ejection direction to lie along the rotation axis of the galaxy, the one direction most closely associated with the ejection of matter from and active galaxy, is equally remarkable. |
This sugsests strongly that the objects have been ejected from NGC 1068. | This suggests strongly that the objects have been ejected from NGC 1068. |
Ilowever. the modest ejection velocities | However, the modest ejection velocities |
GGD 6 (also known as RNO 52), Sh2-235 A (GM1-G6), 235 B (BFS 46 or GM1-G5) and Sh2-235 C (GGD 5 or BFS 47). | GGD 6 (also known as RNO 52), Sh2-235 A (GM1-G6), Sh2-235 B (BFS 46 or GM1-G5) and Sh2-235 C (GGD 5 or BFS 47). |
Sh2-235 A and Sh2-235 B are located about 10’ (5.6 pc) south of Sh2-235, and separated by z40" (0.35 pc) (Blitz,Fich&Stark 1982). | Sh2-235 A and Sh2-235 B are located about $10\arcmin$ (5.6 pc) south of Sh2-235, and separated by $\approx40\arcsec$ (0.35 pc) \citep{Blitz82}. |
. Sh2-235 A is a compact H II region of e20" (0.17 pc) in diameter (Fellietal.1997). | Sh2-235 A is a compact H II region of $\approx20\arcsec$ (0.17 pc) in diameter \citep{Felli97}. |
. For Boleyetal.(2009) the exciting star of Sh2-235 is an early-type Herbig Be star of spectral type B1V. HodappB(1994) found a cluster in Sh2-235 B, which was confirmed by &Felli (1997),, Allenetal.(2005) and Kirsanova (2008). | For \citet{Boley09} the exciting star of Sh2-235 B is an early-type Herbig Be star of spectral type B1V. \citet{Hodapp94} found a cluster in Sh2-235 B, which was confirmed by \citet{Wang97}, , \citet{Allen05} and \citet{Kirsanova08}. |
. Sh2-235 C is a small H II region located about 3.5’ south of Sh2-235 B (Fellietal.2004). | Sh2-235 C is a small H II region located about $3.5\arcmin$ south of Sh2-235 B \citep{Felli04}. |
. It coincides with an optical nebula and is excited by a B0.5 star, presenting a Herbig-Haro object and a partial shell morphology (Fellietal. 1997). | It coincides with an optical nebula and is excited by a B0.5 star, presenting a Herbig-Haro object and a partial shell morphology \citep{Felli97}. |
. Bicaetal.(2003) identify other three ECs in these nebulosities, BDSB 71, 72 and 73. | \citet{Bica03} identify other three ECs in these nebulosities, BDSB 71, 72 and 73. |
Based on the spatial distribution of clusters and kinematics of molecular gas Kirsanovaetal.(2008) point out, that the expansion of Sh2-235 would be responsible for cluster formation in that area. | Based on the spatial distribution of clusters and kinematics of molecular gas \citet{Kirsanova08} point out, that the expansion of Sh2-235 would be responsible for cluster formation in that area. |
They suggest that sequential star formation is triggered by a combination between compression of pre-existing dense clumps by the shock wave, and the scenario. | They suggest that sequential star formation is triggered by a combination between compression of pre-existing dense clumps by the shock wave, and the scenario. |
However, the clusters in Sh2-235 A, B and C appear to be embedded in primordial gas, and star formation cannot be triggered by the expansion of the Sh2-235 ionisation front 1983). | However, the clusters in Sh2-235 A, B and C appear to be embedded in primordial gas, and star formation cannot be triggered by the expansion of the Sh2-235 ionisation front \citep[][]{Lafon83}. |
. On the other hand, Tokunaga&Thompson(1979) suggest that the linear disposition of these star forming regions may be the result of scenario. | On the other hand, \citet{Tokunaga79} suggest that the linear disposition of these star forming regions may be the result of scenario. |
Sh2-232 is an extended H II region with ~40’ diameter (Fig. 2)), | Sh2-232 is an extended H II region with $\approx40\arcmin$ diameter (Fig. \ref{fig:02}) ), |
excited by B stars. | excited by B stars. |
Hodapp(1994) observed two nebulae in Sh2-233 and noticed a young infrared EC, as well as a probable Herbig-Haro object. | \citet{Hodapp94} observed two nebulae in Sh2-233 and noticed a young infrared EC, as well as a probable Herbig-Haro object. |
Porrasetal.(2000) identified two clusters in this region, PCS 2, located around the IRAS 05358+3543 source and Sh2-233 SE ata separation of 1’ (0.5 pc) from each other for an adopted distance to the Sun of 1.8 kpc. | \citet{Porras00} identified two clusters in this region, PCS 2, located around the IRAS $05358+3543$ source and Sh2-233 SE ata separation of $1\arcmin$ (0.5 pc) from each other for an adopted distance to the Sun of 1.8 kpc. |
However, Chan&Fich(1995) estimated a distance of 2.30.7 kpc, based on spectral types | However, \citet{Chan95} estimated a distance of $2.3\pm0.7$ kpc, based on spectral types |
raditional feedback and au X-ray background. | traditional feedback and an X-ray background. |
We note rowever that our barvon-conversion cfiiciency rendus well above the observationallv-derived value: though this xoblem is hardly unique to the present work it reais roublesome. | We note however that our baryon-conversion efficiency remains well above the observationally-derived value; though this problem is hardly unique to the present work it remains troublesome. |
Ou the other laud. less star forination would leave more gas available for AGN accretion. which would likely euliance the relative effect of electromagnetic ecdhack. | On the other hand, less star formation would leave more gas available for AGN accretion, which would likely enhance the relative effect of electromagnetic feedback. |
The enhanced accretion and associated feedback in the radiation-pressure model can also sustain a half-decade ∙∙Star-formation rate of−JPO the hostxvas galaxy NASA eraut for theseveral Cr after a major↿ merger event. although2 the gas required for this extra feedback leads to more star formationin the central regions. which iu turn cau ead to enbhauced ceutral couceutration. | The enhanced accretion and associated feedback in the radiation-pressure model can also sustain a half-decade reduction in the star-formation rate of the host galaxy for several Gyr after a major merger event, although the gas required for this extra feedback leads to more star formation in the central regions, which in turn can lead to enhanced central concentration. |
The ACNX-ray produceslp a significant mass ofLADLE2. virialized. eas⋅⋅⋅ at the present. which tle- X- ⋅:background does not have (vhen ordinary ACN Name hermal feedback. is. preseut: in5 we found sienificant⋅⋅⋅ mass in: hof. dense gas for. a model with. X-:rayav backgroundbackerL but no AGN feedfeedback).dk) whichnich increasesiucreases oth the X-ray ldunuüuositv aud the X-rav half-lieht radius. | The AGN X-ray feedback also produces a significant mass of virialized, soft-X-ray-emitting gas at the present, which the X-ray background does not have (when ordinary AGN thermal feedback is present; in we found a significant mass in hot, dense gas for a model with X-ray background but no AGN feedback), which increases both the X-ray luminosity and the X-ray half-light radius. |
In a sereudipitous final result. wo find that lis ταν feedback is also much more effective than an N-rav backeround iu suppressing αμα. galaxies aud luS flatteniug.. the lowv-1unss slope of the⋅↽ ealaxy Wass its;]⋅ feedback's local origin"n" makingCS↜∙ it‘ o αμ...↴ effective at heating eas some 2003N Cor earlier than the vackerouncl. | In a serendipitous final result, we find that this X-ray feedback is also much more effective than an X-ray background in suppressing small galaxies and thus flattening the low-mass slope of the galaxy mass spectrum, due to its feedback's local origin making it effective at heating gas some $2-3$ Gyr earlier than the background. |
Since ACN are known to cuit N-ravs through their jiost galaxies. we view this study as a vital first step oward a more complete iuodel of AGN feedback. | Since AGN are known to emit X-rays through their host galaxies, we view this study as a vital first step toward a more complete model of AGN feedback. |
Moreover. suce the N-rav luminosity of ACN is relatively well eoustraüned by observations (though not without intrinsic scatter). there is little need for a new free paraueter to joiu the current ay aud er (modulo the effects of iietalliity and dust). | Moreover, since the X-ray luminosity of AGN is relatively well constrained by observations (though not without intrinsic scatter), there is little need for a new free parameter to join the current $\alpha_B$ and $\epsilon_T$ (modulo the effects of metallicity and dust). |
Thus we hope that this “new feedback mode will be eiiploved. iu future SPII simulations of AGN. since it is both undeniably present and. as we lave shown. substautial iu effect. | Thus we hope that this “new” feedback mode will be employed in future SPH simulations of AGN, since it is both undeniably present and, as we have shown, substantial in effect. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅ | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅S | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅St | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Str | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Stru | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Struc | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Struct | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Structu | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Structur | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
supported by NSEaud exaut ASTPID 07-07505 aud reduction iu NNOSATBIC TN ackuowledec support bv the DFC cluster of. excellence “Origin↽⋅⋅ and of. ⇁⋅the Uulverse'. ⋅Structure | JPO was supported by NSF grant AST 07-07505 and NASA grant NNX08AH31G. TN and PHJ acknowledge support by the DFG cluster of excellence `Origin and Structure of the Universe'. |
observations of cirrus iu our own galaxy. | observations of cirrus in our own galaxy. |
Other work ou radiative trausfer modehue of galaxies has been presented by Bianchi et al (1996). Nilouris et al. ( | Other work on radiative transfer modeling of galaxies has been presented by Bianchi et al (1996), Xilouris et al. ( |
1999) aud Popescu et al (2000). | 1999) and Popescu et al (2000). |
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