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Blazar 3C 454.3 in Outburst and Quiescence During 2005-2007: Two Variable Synchrotron Emission Peaks

We monitored the flaring blazar 3C 454.3 during 2005 June-July with the Spitzer Infrared Spectrograph (IRS: 15 epochs), Infrared Array Camera (IRAC: 12 epochs) and Multiband Imaging Photometer (MIPS: 2 epochs). We also made Spitzer IRS, IRAC, and MIPS observations from 2006 December-2007 January when the source was in a low state, the latter simultaneous with a single Chandra X-ray observation. In addition, we present optical and sub-mm monitoring data. The 2005-2007 period saw 3 major outbursts. We present evidence that the radio-optical SED actually consists of two variable synchrotron peaks, the primary at IR and the secondary at sub-mm wavelengths. The lag between the optical and sub-mm outbursts may indicate that these two peaks arise from two distinct regions along the jet separated by a distance of 0.07-5 pc. The flux at 5-35 microns varied by a factor of 40 and the IR peak varied in frequency from<1E13 Hz to 4E13 Hz between the highest and lowest states in 2005 and 2006, respectively. Variability was well correlated across the mid-IR band, with no measurable lag. Flares that doubled in flux occurred on a time scale of 3 days. The IR SED peak moved to higher frequency as a flare brightened, then returned to lower frequency as it decayed. The fractional variability amplitude increased with frequency, which we attribute to decreasing synchrotron-self absorption optical depth. Mid-IR flares may signal the re-energization of a shock that runs into inhomogeneities along the pre-existing jet or in the external medium. The synchrotron peak frequencies during each major outburst may depend upon both the distance from the jet apex and the physical conditions in the shocks. Variation of the Doppler parameter along a curved or helical jet is another possibility. Frequency variability of the IR synchrotron peak may have important consequences for the interpretation of the blazar sequence, and the presence of a secondary peak may give insight into jet structure.


INTRODUCTION
In 2004 August -2005 September, the blazar 3C 454.3 (z=0.859) underwent the largest radio-optical flare in its recorded history. In 2005 May, it briefly surpassed 3C 273 as the optically brightest quasar in the sky in spite of its much greater distance. This flaring event was intensively monitored at all frequencies by observers all over the world, using both ground and space-based observatories (Villata et al. 2006;Pian et al. 2006;Fuhrmann et al. 2006;Villata et al. 2007;Raiteri et al. 2007). The source has since flared twice again with smaller amplitude, offering continuing opportunities to study this spectacular phenomenon, e.g., Villata et al. (2009);Hagen-Thorn et al. (2009) ;Vercellone et al. (2009a); Raiteri et al. (2008a,b).
Blazars, including flat-spectrum radio-loud quasars and BL Lac objects, are characterized by strong variability and high polarization. According to active galaxy unification models, they are viewed within ∼ 20 • of the radio jet axis. The spectral energy distribution (SED) is dominated by relativistically boosted emission from the core of the jet. The core is strongly variable at all frequencies, owing to rapid movement and changes in the jet on sub-parsec scales. Very-long baseline interferometry typically shows jet components with apparently superlu-ogle@ipac.caltech.edu minal motion at milliarcsecond resolution, demonstrating the importance of relativistic effects.
The SEDs of blazars are typically characterized by two major bumps, one which peaks in the radio through Xray bands, and the other which peaks in X-γ ray bands. The low frequency bump is highly polarized, indicating synchrotron emission from relativistic electrons in a magnetized jet. The high frequency bump is thought to be produced by inverse-Compton scattering of photons by relativistic electrons. However, the source of the photons is a subject of debate, ranging from the synchrotron photons in the jet itself to external photon fields (e.g., Wehrle et al. 1998). In some blazars, a significant fraction of the optical-UV thermal emission as well as broad optical-UV emission lines may arise from an accretion disk (Smith et al. 1988). See, for example, Ghisellini & Tavecchio (2009) and references therein, for a review of models and their application to the class of flat spectrum radio quasars.
Theoretical unifying schemes for γ-ray bright blazars have been proposed: the more luminous the blazar, the lower the peak frequencies of the synchrotron and Compton bumps Ghisellini et al. 1998). Physically, the down-shift in peak frequency can be caused by increased Compton cooling at higher radiation densities. According to the external-Compton (EC) model, photon fields from the accretion disk, broad-line region (BLR), and dusty torus act to sap energy from the relativistic jet. Hence, luminous flat-spectrum radio quasar (FSRQ) SEDs peak (on average) at lower frequencies than do those of BL Lac objects. Similarly, we might expect accretion disk luminosity variations in a single object to be accompanied by changes in the frequency of the synchrotron peak.
However, Compton cooling may be counterbalanced by re-acceleration of high energy electrons in jet shocks. As a disturbance travels down the jet, it may run into inhomogeneities internal to or external to the jet. At such locations, the bulk relativistic kinetic energy and energy in the electromagnetic field is dissipated, heating the electrons, which emit copious synchrotron photons as they gyrate about the magnetic field lines. The increased number of high energy electrons will shift the synchrotron peak to higher frequencies. Thus the peak synchrotron frequency is determined by a balance between shock heating and Compton cooling, and should vary as jet shock components evolve. An exact correspondence between accretion luminosity and peak synchrotron frequency is therefore not expected.
The IR is a crucial band since it covers the primary synchrotron emission peak of FSRQs, which may provide seed photons for the Compton γ-ray bump. In addition, while the base of the jet is optically thick at radio frequencies, it is optically thin at mid-IR through UV frequencies. Thus IR observations are less subject to optical depth effects than radio observations and can be used to probe short variability time scales which arise close to the origin of the jet. If flaring episodes are caused by shocks (Marscher & Gear 1985), then flares should occur simultaneously across the IR through UV wavebands (Lainela 1994;Savolainen et al. 2002). Therefore, IR spectral variability should reflect changes in the energy distribution of relativistic electrons. Variability on day to month time scales can be used to study the evolution of the shock as it propagates along the jet and through the environment of the active galactic nucleus.
The blazar 3C 454.3 has undergone a number of flaring episodes, with bright radio outbursts occurring roughly once every 6 years (Ciaramella et al. 2004).
Historically, the synchrotron emission from 3C 454.3 typically peaks in νF ν in the far-IR (∼ 100 µm). IR emission observed by ISO had a flux of 34 mJy at 12.8 µm (Haas et al. 2004). Excess emission in the SED at 60 µm was attributed to thermal emission from 100 K dust by these authors, but we will present evidence that this emission is highly variable. The X-ray spectrum observed with Beppo-SAX was characterized by an 0.5-200 keV power law with spectral index α = −2.3 (Tavecchio et al. 2002).The greatest amount of energy is emitted in the Compton bump, which peaks at ∼ 10 MeV. Gamma-ray emission was observed at energies of 0.05-4 GeV by the EGRET and OSSE instruments on the Compton Gamma-Ray Observatory (Hartman et al. 1993;McNaron-Brown et al. 1995;Zhang et al. 2005). Gamma-ray detections after the data in this paper were obtained include those by AGILE (Vercellone et al. 2009a) in 2007, by the LAT instrument on Fermi in 2008 (Abdo et al. 2009), and by AGILE in 2008(Donnarumma et al. 2009a); see also the review by Vercellone et al. (2009b).
We were awarded director's discretionary time to observe 3C 454.3 daily with the Spitzer Infrared Spectrograph (IRS), the Infrared Array Camera (IRAC), and the Multiband Imaging Photometer (MIPS) over a period of 4 weeks in 2005 July, during normally scheduled instrument campaigns. We also took coordinated Spitzer MIPS, IRS, IRAC and Chandra High Energy Transmission Grating (HETG) observations in 2006 December -2007 January. We obtained observations from long term optical monitoring programs at Colgate University and Palomar Observatory and from calibrator monitoring at the Submillimeter Array (SMA). In this paper, we present these data and discuss variability and the nature of synchrotron emission flares from the blazar 3C 454.3. HISTORY OF 3C 454.3 IN 2005 Observations that were made during the same time period as those reported in this paper included several large multiwavelength campaigns, summarized here. Following the alert of an optical flare in 3C 454.3 by Balonek (2005a) and Balonek (2005b), Fuhrmann et al. (2006) obtained near-IR/optical photometry starting two days later. Villata et al. (2006) made the first report of the WEBT multiwavelength campaigns, including ToO pointings by Chandra and Integral, that followed the discovery of the May 2005 optical flare. Raiteri et al. (2007) reported the detection of the "small and big blue bumps" during the low state in 2006. Villata et al. (2007) derived the delay between centimeter-band radio and optical light curves during [2005][2006]. Giommi et al. (2006) observed 3C 454.3 with the robotic Rapid Eye Mount optical/near-infrared telescope and with the Swift satellite in April-May 2005. Raiteri et al. (2008a) presented multifrequency observations by the WEBT andXMM-Newton in 2007-2008, overlapping with AGILE's November 2007 observations. They also used 1.3 mm monitoring data from the SMA, some of which is also used in our paper, over the time period 2005-2008 (see their Figure 4) as well as 1.3 mm data from Pico Veleta, Spain. X-ray data from Swift, Chandra, XMM-Newton, and Integral data from the 2005-2007 period have been presented by Giommi et al. (2006), Villata et al. (2006), Raiteri et al. (2007) and Pian et al. (2006). We describe the results of these papers as follows.

A BRIEF
1. Fuhrmann et al. (2006) observed 3C 454.3 with the Automatic Imaging Telescope of Perugia University Observatory and the Rapid Eye Mount telescope in Chile, obtaining V,R,I and H band photometry. The spectral index over these bands showed no strong significant changes.
2. Villata et al. (2006) presented results up through September 2005, finding that the source was redder when brighter. A mm outburst occurred in June-July 2005, followed months later by the 37-43 GHz peak. Chan-dra and Integral X-ray observations in May 2005 showed unusually high fluxes.
3. Giommi et al. (2006) observed 3C 454.3 with the robotic Rapid Eye Mount optical/near-infrared telescope and with the Swift satellite in April-May 2005. They found that the optical and ultraviolet flux doubled within a single 1.5 × 10 5 second exposure, the XRT (2-10 keV) X-ray flux varied little during the same time. However, on time scales of a few days, the BAT (15-150 keV) X-ray flux varied by more than a factor of three; in contrast, the average level of the optical-ultraviolet flux was approximately constant between the two UVOT observations of 24 April and 17 May 2005.
4. Villata et al. (2007) reported multiwavelength observations during [2005][2006]. Their analysis suggested that the big radio flare (43-37 GHz) in early 2006 was associated with a minor optical flare in October-November 2005, not with the spring 2005 major optical flare. A combination of disturbances traveling down the jet and changes of viewing angles of different emitting regions, with concomitant changes in Doppler boosting, were found to explain the radio delays with respect to the optical emission.
5. Raiteri et al. (2007) detected the "small and big blue bumps" during the low state in 2006. Observations from K-band through ultraviolet showed a small peak in optical band (V), an upturn in ultraviolet from Uband to UVW1 and UVM2 and an upturn at the infrared end. The infrared end was probably variable synchrotron emission, complicated by a non-variable H-alpha line (reported by Raiteri et al. (2008a) via spectroscopy in J band). The "small blue bump" is probably a blend of iron lines, Mg II lines and Balmer continuum from the broad line region. The ultraviolet upturn, interpreted as the beginning of the big blue bump was probably from thermal emission of the accretion disk. All these physical emission mechanisms contributed during various activity levels. The underlying accretion disk and broad line contributions were visible most clearly when the synchrotron emission was at low ebb. The X-ray spectra could be fitted with a power law, but seemed to require extra N H absorption some of the time. Alternatively, there may have been spectral changes in the soft X-ray range.
6. Raiteri et al. (2008a) presented multifrequency observations by the WEBT and XMM-Newton in 2007-2008, overlapping with AGILE's November 2007 observations. They also used 1.3 mm monitoring data from the SMA, some of which is also used in our paper, over the time period 2005-2008 (see their Figure 4) as well as 1.3 mm data from Pico Veleta, Spain. Correlation of the WEBT R-band optical data with the SMA and Pico Veleta data showed that the millimeter fluxes lagged the optical fluxes by 40-80 days, with Discrete Correlation Function maximum signal at 65 days. Taking only nonoutburst data for the correlation yielded a DCF maximum at 20 days. They suggested that the "jet regions emitting the optical and mm radiations are now better aligned than in the past, and/or that the opacity in the jet has decreased, allowing the release of mm radiation closer to the optically emitting zone." We return to this correlation later in our paper. Several multiwavelength SEDs during 2005-2007 are shown and analyzed. From the X-ray-optical correlation, they found that the X-ray, near-infrared and optical emission would be produced in the same spatial region.

Submillimeter Array Observations
The Submillimeter Array (SMA) is an 8-element radio interferometer located atop Mauna Kea, Hawaii, which operates in the 1.3 mm, 850 µm, and 450 µm atmospheric windows (Ho et al. 2004). In the 1.3 mm and 850 µm windows, quasars are utilized as amplitude and phase gain calibrator sources during routine observing, and the flux density scale is derived through observations of standards in each session, typically solar system sources (with Uranus, Neptune, Titan, Ganymede and Callisto being the most used). The SMA supports a program to calibrate quasar flux densities using both routine science and dedicated calibration observations, and provide their histories to SMA users and the wider astronomical community. The Submillimeter Calibrator List, containing the flux measurement histories of several hundred quasars, can be found in the Tools section of the SMA Observer Center 1 , with further details provided by Gurwell et al. (2007). Data for 2005-2007 are presented in Table 1 and Figure 1 for the 850 µm and 1 mm bands.

Submillimeter Flares
The sub-mm data, shown in Figure 1a Figure 1b shows the extraordinary flare in 2005 and a surprisingly large amount of substructure during the 2-month period around the peak, including five dates on which the flux exceeded 40 Jy. We think the substructure is real rather than random noise because the variability correlates well between the 230 and 345 GHz band measurements with several measurements in a row showing consistent increasing or decreasing flux density. We can not completely exclude systematic error, for example, observer bias in knowing what the preceding measurement value was.
The SMA is sensitive to only a single polarization which also rotates on the sky as a function of elevation. Anecdotal reports by other observers indicated that 3C 454.3 was not polarized more than a couple of percent. Given the complex rotation relative to the source, it is unlikely that any significant structure in the time variability would be related to changes in the polarization angle, but it cannot be excluded without considerable extra work beyond the scope of this paper.
For about 98% of the measurements, the frequency quoted is in fact the local oscillator (LO) frequency, and the flux measurement is the average of the lower and upper sideband values (in the heterodyne mixing, the sky frequencies are mixed down to the more manageable IF frequency range of 4-6 GHz, but we accept both sidebands, e.g. LO +/-(4-6 GHz). The use of interferometry and the complex correlator allows us to unambiguously separate the sidebands, so we really get two flux measurements, separated by 10 GHz. The SMA 3C 454.3 data yield a spectral index that typically lies near ν −0.5 , thus we would expect only a 2% drop in flux from 220 to 230 GHz (1.3 mm band), and only 1.5% from 335 to 345 GHz (850 µm band). The peak substructure in the sub-mm flare is definitely not caused simply by changes in the frequency of observation. We observed 3C 454.3 with Spitzer IRS, IRAC, and  MIPS at several epochs during 2005-2007 (Tables 2 and  3). Note that the three instruments were operated separately (according to their regularly scheduled campaigns of typical length 1-2 weeks) so we do not have simultaneous coverage in the respective IR bands. First, we observed the source once per day from 2005 June 30 -2005 July 27 to track mid-IR spectral variability following the 2005 May outburst. We conducted a series of 14 daily IRS observations between 2005 June 30 and July 12, followed by 12 daily IRAC observations between July 14 and July 26, followed by two MIPS observations on July 27, the last day of the Spitzer visibility window. We observed the source again on 2006 Dec 20 -2007 Jan 1, the first of two coordinated Spitzer and Chandra observation sequences (the second is reported by Wehrle et al. 2010, in preparation Tables 2 and  3.

IRAC
The IRAC images (Fig. 2) on 2005 July 14 (bright state) and 2006 Dec 25 (faint state) show the unresolved quasar in a field of stars and foreground galaxies. The Spitzer Science Center (SSC) pipeline mosaic data (version 14.0.0) were used for aperture photometry carried out with SSC APEX software package in beta release which was the only Spitzer software available for photometry when we began this work. The apertures used were 2, 4, 5, and 5 pixels in radius (1 pixel = 1.2 arcsecond) at 3.6, 4.5, 5.8 and 8.0 µm, and corrected using the values given in Table 5.7, p. 53 of the IRAC Data Handbook, Version 3.0. We cross-checked the computed quasar fluxes for a range of apertures with those of a nearby star of roughly comparable brightness; the errors are consistent with the 1.5% systematic errors plus 3% absolute calibration uncertainty of the IRAC (Reach et al. 2005). Despite the unpredictable brightness of the source, none of the IRAC images were saturated. Following the release of a new version of IPAC's Skyview software (courtesy of B. Hartley (IPAC)), we used it to reduce the images independently. The APEX and Skyview results agree to within two percent, except on 2005 July 14, where the APEX value exceeds the Skyview value by 10% at 3.6 µm. The Skyview fluxes at 3.6, 4.5, 5.8, and 8 µm are listed in Table 4.
Serendipitously, during the faint state, the 3.6 and 4.5 µm IRAC images displayed a faint, 10-σ feature about 6 " from the quasar at PA ∼ 150 degrees. IRAC detections have been made of arcsecond-scale jets from other blazars, however, this feature appeared on the opposite side from the 8" radio jet at PA -45 degrees (e.g., Cooper et al. 2007). The feature was present in all individual frames when the source was faint, including those from another of our Spitzer programs. It was not associated with typical IRAC instrumental artifacts (J. Surace, private communication). Working with Marco Chiaberge, we pulled archival ACS Hubble images (F. Tavecchio, PI) of the source during a faint episode in August 2004, and found a plausible identification of the IRAC feature with two galaxies, conflated by the much larger IRAC PSF. The galaxies' brightness and colors indicated that they could be located at roughly the same distance as 3C 454.3 (W.C. Keel, private communication). Figure 3 shows the IRAC and Hubble images side by side.

MIPS
The MIPS images (Fig. 4), taken when the source was bright on 2005 July 27 at 01:48 UT and 15:11 UT, and when the source was faint on 2007 January 1, show the unresolved quasar. In the full 24 µm images, the bright quasar and a few faint galaxies were visible. At 70 µm, the quasar was isolated on fairly smooth background. The first Airy ring is clearly visible in each image. In the 2005 July 27 datasets, evidence of "soft saturation" was found in the 160 µm data (the pipeline images showed a doughnut-like structure instead of a smooth peak). Special MIPS custom processing on the BCD-level data was performed to obtain useful fluxes by MIPS Instrument Team member K. Gordon (see Gordon et al. (2005) for software description). For 24 and 70 µm data, the pipeline mosaic data (version 13.2.0) was adequate for our use. The 2007 Jan 01 data were not affected by soft saturation; we used pipeline data (version 13.2.0). The IPAC Skyview analysis package was used for aperture photometry. At 24 µm, we used a 13" radius aperture and applied an aperture correction of 1.164 (Table 3.13, p. 33 of MIPS Data Handbook version 3.2). At 70 µm, we used 35" radius aperture and applied an aperture correction for a ν −2 spectral index of 1.197 (Table 3.14, p. 33, MIPS Data Handbook version 3.2). For our 160 µm data on 2007 Jan 1, we used 50" radius aperture and applied an aperture correction of 1.470 (Table 3. 16,p. 35,MIPS Data Handbook v. 3.2).
Additional confidence in the Spitzer low state 160 micron flux density of 0.22 Jy with 30% systematic uncertainty, measured on 2007 Jan 1, is provided by the completely independent Infrared Space Observatory Photometer measurements while the quasar was in an earlier low state: ISO-Phot measured the 120 and 180 micron flux densities of 0.187 and 0.248 Jy on 1997 December 18 (Haas et al. 2004). ISO-PHOT and Spitzer MIPS were calibrated using different primary and secondary calibrators, including asteroids. The 3C454.3 maps made by ISO and Spitzer were in different orientations on the sky.

IRS Low Resolution
IRS was used in standard staring mode to take lowresolution spectra with the Short-Low (SL2,1) and Long-Low (LL2,1) modules (Figs. 5-6). Because of the unex-pectedly bright flux level, digital saturation occurred in some of the 2005 July exposure ramps. This effect is automatically corrected for in the data reduction pipeline by throwing out the saturated samples, leading to lower effective exposure time for the brightest pixels and a corresponding reduction in S/N. Several pixels in the LL1 order of the epoch 14 spectrum had only 2 unsaturated samples in each ramp, leading to the noise spike at 22 µm (Fig. 5). To guard against this problem, we reduced the exposure times and increased the number of exposures in subsequent epochs.
We have devised a fringe-correction algorithm to ameliorate spectral fringes in the LL spectra. The wavelengths of the detector fringes depend on where exactly the source is placed in the slit, which varies with the pointing error of the telescope. The detector fringe patterns were extracted from flat field images and used to derive fringe correction functions. The fringe spectrum was shifted and divided by itself to generate a correction curve to minimize the residual fringes in the source spectrum. This procedure reduced the residual fringe amplitudes from < 4% to < 2% in LL.
We found significant offsets between low-resolution spectral orders, which were likely caused by random pointing offsets. This effect was most pronounced in the SL1 order because of the small size of the slit relative to the telescope point-spread function. Multiplicative corrections of 1.0-1.2 (constant with wavelength) were applied to match the flux in each of the SL2, SL1, and LL2 spectral orders to LL1 to within an accuracy of 1%. The corrections were largest in the epoch 6-8 SL1 spectra (8-20%). For these epochs, there was a residual convex downward curvature of ∼ 3% in SL1 caused by wavelength dependent slit losses, which we did not correct, but which do not affect our conclusions.
We used flux calibrations provided by the SSC, employing low-order polynomial fits to standard star observations to convert from electron s −1 to Jy. The absolute flux calibration accuracy is limited by the uncertainty of the stellar atmosphere model for the flux calibration standard HR 7341, which is estimated to be 3 − 5% across the IRS wavelength band (Decin & Eriksson 2007;Decin et al. 2004). The relative flux accuracy of each epoch is limited by slit losses. Repeatability observations of the standard star HD 173511 give independent 1-σ flux calibration uncertainties of 2% for SL2, 1% for SL1 and LL2, and 3% for LL1. After matching all orders to LL1, the relative flux accuracy is 3%. We achieve S/N values of ∼ 60−120 in the low-resolution spectra, limited by remaining systematic uncertainties in the flat field at a level of 1 − 2%.
There are no emission lines or absorption lines stronger than 2% of the continuum flux in the 2005 June-July low resolution spectra (Fig. 5). There is also no evidence for any significant broad silicate absorption or emission or any emission from polycyclic aromatic hydrocarbons (PAHs). Small deviations of < 2% come from uncertainties in the flux calibration introduced by residual detector fringing effects. The most prominent instrumental feature is a ∼ 10% bump at 14 µm (the so-called SL1 "teardrop"), which may be caused by scattered light in the detector. There is a similar instrumental bump at 7 µm and a dip at ∼ 5.5 µm, both in SL2 and of unknown origin.

IRS High Resolution
We also took high resolution spectra ( Fig. 7) with the Short-High (SH) and Long-High (LH) modules at the beginning and end of the 2005 IRS campaign (June 30 and July 13). These observations were designed to search for any narrow emission or absorption features in the mid-IR spectrum. The high resolution spectra are divided by a factor of 1.04 to match the low-resolution flux calibration, since there is a level mismatch between the two flux standards used in the S15 pipeline to calibrate low and high-resolution data 2 . No sky subtraction was performed. The sky background was small compared to the source flux in SH for the brighter of the two epochs. However, it became significant for the larger LH slit and for the fainter epoch.
We subtracted a second-order polynomial fit to each order of the high-resolution spectra in order to get a closer look at any possible emission line features (Fig. 8). On inspection, it is clear that there is no significant detection at either epoch of any of the following quasar or host galaxy emission features at a redshift of z = 0.859: There is a possible detection of [Ne vi] λ7.65 µm at 14.22 µm. However, it is most likely caused an instrumental defect or noise at the edge of SH order 14.

Comparison of 2005 June-July with 2006 December IRS Observations
The 2005 June-July Spitzer spectra (Fig. 5) were taken more than two months after the peak of the 2005 optical outburst. Even so late after the peak, 3C 454.3 was a factor of ∼ 20 brighter at mid-IR wavelengths than at any previously published epoch. The peak flux during this period reached 1.14 ±0.03 Jy at 12 µm and 2.46 ±0.07 Jy at 24 µm (Table 10). The mean spectrum was fit by a power law F ν ∼ ν −1.2 at 6-12 µm, with an exponential cutoff at longer wavelengths. The primary synchrotron bump peaked at an observed frequency of 1-4×10 13 Hz (30-7 µm).
In 2006 December, we found 3C 454.3 in a relatively low state (Fig. 6). The low state spectrum is well fit by a simple power law with spectral index α = 1.3. We find no evidence of emission lines or silicate emission stronger than 3% of the continuum. This suggests that the lowstate IR SED is dominated by synchrotron emission from the jet. The flux was 28 ±1 mJy at 12 µm, a factor of 40 down from the maximum in 2005 July and slightly fainter than the historical ISO measurement of 34 mJy at 12.8 µm (Haas et al. 2004).

Optical Observations
Since 1988, T. Balonek has conducted blazar monitoring observations with Colgate University's Foggy Bottom Observatory. As described by Kartaltepe & Balonek (2007), observations are made "using a Ferson 16 inch Newtonian/Cassegrain reflecting telescope equipped with a Star I CCD system. The images range in exposure time from 1 to 5 minutes (most images being 2 minutes) and include the B, V, R, and I filters, designed by Beckert (1989) to conform to the Johnson-Cousins system...The images were reduced using standard IRAF packages and customized scripts written to facilitate the data handling for our system." In the case of 3C 454.3, V, R and I band data were obtained. The photometry for all of the images was calculated using the IRAF apphot package with a 16" diameter aperture and a sky annulus with an inner diameter of 20" and an outer diameter of 40". Some of the Colgate data were previously shown by Raiteri et al. (2008a). The Foggy Bottom Observatory data, uncorrected for absorption, are listed in Tables  6-8.
Since 2005, the members of the NASA Space Interferometry Mission Key Project on Active Galactic Nuclei, led by A. Wehrle, have conducted blazar monitoring observations using the Caltech automated 1.5m telescope at Palomar Observatory (Cenko et al. 2006). We intensified monitoring of 3C 454.3 when it flared in 2005. Data in filter bands V and R were obtained 3 . Data were reduced with IRAF using aperture photometry, with an 4.56" diameter aperture and a sky annulus of 3.8" radius and 3.8" width. The comparison star sequence is given online 4 and originally by Craine (1977), Angione (1971), Fiorucci et al. (1998), Smith & Balonek (1998), and Raiteri et al. (1998). We used Star H, with V=13.64 and R= 13.10, as the primary comparison star. The Palomar Observatory data, uncorrected for absorption, are given in Tables 9 and 10. 3.4. Chandra Observations Chandra High Energy Transmission Grating (HETG) observations were taken simultaneously with our Spitzer MIPS observations, for 2.145 ks from 21:35-22:42 UT on 2007 January 1. The grating was used to reduce potential photon pile up effects on the ACIS-S detector. The data were reduced using standard point-source grating analysis threads under CIAO 4.0. The MEG and HEG ±1 order grating spectra were grouped to have a minimum of 20 counts/bin. We fixed the hydrogen column density for model fits at N H = 6.5×10 20 cm −2 (Dickey & Lockman 1990) and assumed standard elemental abundances to estimate Galactic absorption. The best fit to the X-ray continuum was obtained with a single power law of slope Γ = 1.53 ± 0.12 and flux at 2-10 keV of 5.8 ± 0.2 × 10 −12 erg cm −2 s −1 (8.14 × 10 −4 ph s −1 cm −2 ). Adding absorption intrinsic to the host galaxy to the model did not improve the fit. Similar fits were obtained using either ξ 2 or C-statistics. We report fit parameters using the latter, because of small photon number statistics in spite of grouping the data.

Optical Light Curves
The optical light curves are shown in Figure 9. We scaled the Colgate V band data by 1.15 to match the Palomar V data in constructing the light curves; the discrepancy is probably due to slight mismatches in filters.

Infrared Light Curves
The 2005 June-July infrared light curves for IRS and IRAC are shown together with the optical light curves in Figure 9. The IRS portions of these curves are shown in more detail in Fig. 10ab. We compute mean 6-12, 6-24, and 12-24 µm spectral indices to characterize the spectral hardness Fig. 10c. The photometric variability is similar in all wavebands, with greatest amplitude at the shortest wavelengths. The 6 µm flux increases by a factor of 2 during 2005 July from 240 ± 20 mJy at epoch 3, flaring to 470 ± 30 mJy at epoch 14. All of the dips and peaks in the light curves occur simultaneously at 6, 12, and 24 µm, to within 0.5 day. There is no measurable lag.
Significant variations are seen in the spectral index as the flux changes in 2005 June -July (Fig. 10). In particular, as the flux increases from epoch 4-5, the slope becomes softer. The flux and spectral slope show only small changes from epoch 5-10. This period is followed by a large flare from epochs 12-14, when the flux increases by 85% at 6 µm and the spectrum hardens. The flux falls and the spectrum softens again in the epoch after the peak of the flare.
The 6-24 µm spectral index varies from a maximum of 1.41 ± 0.06 at epoch 9 to a minimum of 1.17 ± 0.06 at epoch 13 (Fig. 10). The spectral index uncertainties correspond to the IRS flux repeatability of 2% at 6 µm and 3% at 24 um, as given in Section 2.2.3. There is an anticorrelation between spectral index and 6 µm flux (Fig. 11) during 2005 July. The SED becomes harder at high frequencies as the flux increases. The peak of the SED also moves to higher frequency as the flux increses: from 1.8 − 2.4 × 10 13 Hz to 2.2 − 4.0 × 10 13 Hz. Whenever the 6 µm flux increases (epochs 3-4, 7-8, 9-10, 11-14), the spectral index decreases. This is a consequence of the larger variability amplitude at 6 µm than at 24 µm. The variability amplitude may be moderated by optical depth to synchrotron self-absorption, which increases with wavelength.

Correlation of Optical and Infrared Light Curves
An inspection of the optical through infrared light curves (Fig. 9) reveals correlated variability from the V-band to 24 µm. In particular the epoch 14 flare seen in the IRS light curves (Fig. 10), is seen in all of the optical bands as well. However, a close comparision of the IRAC/IRS 8.0 µm band with the R-band light curve (Fig. 12) shows significant differences in flaring activity 5 . In particular, the 8.0 µm band shows a bright flare at days 15-20, following the epoch 14 flare, which is much less pronounced in the R-band. There are clear differences in the relative amplitudes of several other smaller flares, which change continuously with wavelength between the IR and optical bands. This behavior may indicate multiple flaring components in the 3C 454.3 SED which have peak amplitudes at different wavelengths. (See Section 4.4.) A cross-correlation of the IRAC/IRS 8.0 µm and Rband light curves shows a pronounced peak with correlation amplitude 0.7 at a lag of 0 ± 0.5 days (Fig. 12c). This confirms the overall impression that the flaring activity is correlated across the optical and mid-IR bands. However, the cross-correlation curve is broadened, with a secondary peak at a lag of 4 days. This lag is similar to the time between the two brightest 8.0 µm flares, which may produce an accidental (rather than causal) correlation signal.

Mid-Infrared-Optical SEDs
The mid-infrared to optical spectral energy distributions are shown in Figure 13. From the Figure, we observe the following.
1) During the period 30 June -13 July 2005, there was less variation at IRS bands than in VRI bands, but during the following ten days, 14 July 2005-25 July, there was more variation at IRAC bands than at VRI, including the 5.8 and 8.0 µm bands that overlap in wavelength with the IRS.
2) The SED slope over mid-infrared through nearinfrared to optical bands was negative in both bright (2005) and faint (2007) states. When the source was a factor of 6-10 weaker (in 2006 July and 2006 December) than it was during our July 2005 observations, Raiteri et al. (2007) observed that the underlying nearinfrared-optical (I-R-V bands) SED slope was positive, unlike our 2005 observations. In their data, the optical emission peaked at V, then decreased in Swift B and U bands and rose again at Swift W1 and M2 bands. These data, taken together, are consistent with the Raiteri et al. detection of an underlying "small blue bump" (probably a blend of iron lines, MgII lines and Balmer continuum from the broad line region) whose V-band-peaked emission is visible only when the synchrotron emission is comparatively faint, and whose signature is overwhelmed by synchrotron emission when the source is flaring 6-10 times brighter. Our observations do not cover the Swift bands which show the rise in thermal emission from the accretion disk.
3) The SED does not rise or fall by a constant amount across the mid-infrared to optical bands from day to day. Moreover, the individual band-to-band spectral indices are not the same from day to day. There is no obvious wavelength range over which we see the beginning and ending points of added emission. That implies the newly injected population of leptons giving rise to the synchrotron emission is broad in energy, but we have to keep in mind that the particles radiating at a given frequency are from both old and new populations. Previously injected populations contribute at lower frequencies long after that population's high energy electrons no longer radiate at high frequencies, assuming no change in the ambient magnetic field. 4) Energy gained in a day at optical bands can be lost in a day, but it takes longer to both gain and lose energy at mid-infrared bands. We see no evidence of asymmetric rise or fall in the mid infrared or optical flares.

Comparison of Outburst and Quiescent X-ray
Fluxes and Spectral Indices On 19 May 2005, during outburst and a month earlier than our first Spitzer observations, Chandra observed 3C 454.3 for 112 ksec using the HRC-LETG, as a Target of Opportunity within a program on flaring blazars (F. Nicastro, PI; Villata et al. 2006). The 0.2-8 keV spectrum was fitted with a power law of photon index Γ = 1.477 ± 0.017, with N H = (1.34 ± 0.05) × 10 21 cm −2 , more than twice the Galactic value (a misprint in Villata et al. 2006 has been corrected). During our 1 January 2007 Chandra observations, obtained using the HETG mode when the source was quiescent, the 2-10 keV spectral index was quite similar (Γ = 1.53 ± 0.12) using the Galactic N H value of 6.5 × 10 20 cm −2 . The May 2005 deabsorbed fluxes were F = (5.5 ± 0.2) × 10 −11 erg cm −2 s −1 and (8.4 ± 0.2) × 10 −11 erg cm −2 s −1 in the 0.2-2 keV and 2-8 keV bands (respectively), compared to the January 2007 flux at 2-10 keV of 5.8 ± 0.2 × 10 −12 erg s −1 . Hence, the 2-10 keV X-ray flux dropped by a factor of ∼ 10 while the spectral index remained the same, but the bright state required twice as much N H along the line of sight as the quiescent state. An Integral observation (Pian et al. 2006) made 15-18 May 2005 was fitted with a power law of index Γ = 1.8 ± 0.1 using the same (Dickey & Lockman 1990) Galactic N H value The flux in the band 3-200 keV was 5.45 × 10 −10 erg cm −2 s −1 .

The Spectral Energy Distributions: Two
Synchrotron Peaks We use our multiwavelength photometry and the IRS spectra together in constructing radio-X-ray SEDs (Figure 14); representative data are listed for convenience in Tables 11 and 12. In constructing the SEDs, the photometry was corrected for Galactic absorption at V, R and I bands by 0.355, 0.286, and 0.208 magnitudes, respectively. The Colgate V band data were also scaled by 1.15 to match the Palomar V data, as mentioned above.
During 2005 June 30-July 12 (high state), the primary SED peak fell in the IRS 5-35 µm band (Fig. 14). Strong variability of the continuum level, slope, and peak frequency are characteristic of synchroton emission from a relativistic radio jet. The largest variability amplitudes are seen in the IRAC 3.6 µm and near-IR (I) bands and the smallest amplitude at 35 µm. The turnover and lower variability amplitude at lower frequencies may indicate the onset of optically thick synchrotron emission. Without long-wavelength IR coverage, we lost track of the primary synchrotron peak location during the July 14-26 IRAC campaign. When the MIPS observations picked up on July 27, the IR peak had faded somewhat and shifted to the MIPS 70 µm band.
In the 2006 December 25-2007 January 2 low state, the IR SED peaked in the MIPS 70 µm FIR band. The lowstate SED was similar to that reported by Haas et al. (2004), and to other low-state measurements from NED published over the period 1979-1995 (Fig. 14). One striking difference was the higher 60 µm flux measured by Haas et al. (2004), which was ∼ 4 times greater than our MIPS 70 µm flux, but less than the 2005 peak. This is additional evidence that the 60 µm flux is highly variable, and indicates that there was little if any contribution from thermal dust emission during the Haas et al. (2004) observation. It is particularly evident during our low-state Spitzer observations (Fig. 14) that there are two peaks in the radio-IR SED, as seen earlier in the lowstate observations, made by combining ISO data from 1996 Dec 12 and 1997 Dec 18, of Haas et al. (2004). As we recall from Section 3.2.2, the veracity of the Spitzer 160 micron flux measurement in 2007 January is backed up by the detection of a similarly low level in 1997 by ISO.
During both the low and high states, the primary SED peak occurs at MIR-FIR wavelengths, depending on source brightness. The secondary peak occurs in the sub-mm, near 850 µm. Strong variability of both bumps indicates that they are both composed of synchrotron emission from the relativistic jet. At the time of the 2005 July 27 MIPS observations, two bumps are also apparent, with the minimum falling somewhere between 160 µm and 850 µm. The sharpest contrast between the two peaks is seen during the ISO epochs, when the infrared peak was defined by the 0.706 Jy peak at 60 µm.
In contrast, non-blazar quasars have two thermal peaks in their SEDs: an ultraviolet peak from the accretion disk and an infrared peak due to reprocessing of ultraviolet light by a dusty torus. We have plotted the mean Richards et al. (2006) QSO SED, compiled from SDSS, Spitzer, near-IR, GALEX, VLA, and ROSAT data in Figure 14, redshifted and scaled for comparison to the 3C 454.3 low state. The UV thermal peak is matched to the big and small blue bumps observed by XMM OM (Raiteri et al. 2007). The IR bump in the low-state 3C 454.3 SED is significantly redder and peaks at a much lower frequency than the IR bump in the mean QSO SED. A thermal SED that peaks at 70 µm is difficult to model with a standard dusty torus (e.g., Levenson et al. 2007), and would instead require that the QSO be deeply buried in an optically thick sphere of cold dust, contrary to the lack of any silicate absorption in the MIR spectrum and presence of an unobscured Big Blue Bump (BBB) component in the UV SED.
We can also confidently rule out a significant thermal contribution from cold dust heated by star formation in the host galaxy of 3C 454.3. In Figure 14, we over-plot the SED template constructed for a star formation dominated hyper-ULIRG (Rieke et al. 2009) with bolometric luminosity 1×10 13 L ⊙ , scaled up by a factor of 2.2 to match the low-state 70 µm (isotropic) luminosity of 2.9 × 10 46 erg s −1 . Hyper-ULIRG galaxies are extremely rare in the universe and it would be very unusual for a such a high-luminosity, radio loud, type 1 quasar to have a starburst dominated FIR SED. Furthermore, the low state 5-70 µm power law with spectral index −1.3 ( § 3.2.5) is is much flatter than the spectral index of the template hyper-ULIRG SED. We therefore conclude that the jet provides the dominant contribution to the low-state SED at 70 µm.

Interpretation of the 2 Synchrotron Peaks
The presence of two synchrotron peaks has also been seen or suggested in at least 3 BL Lac objects, including BL Lac (Ravasio et al. 2003;Raiteri et al. 2009), Mrk 501 (Villata & Raiteri 1999), and Mrk 421 (Donnarumma et al. 2009b). The frequencies of the two peaks are in the near-IR to optical and far-UV to X-ray bands for BL Lacs, in contrast to the sub-mm and far-IR for the FSRQ 3C 454.3. Ravasio et al. (2003) suggest four possible explanations for the double peaks in BL Lac: 1) higher than normal ratio of optical to X-ray extinction, 2) bulk Comptonization, 3) Klein-Nishina effect on Compton cooling, and 4) two synchrotron emission regions. We can immediately rule out the first two explanations for the two peaks in 3C 454.3 because they occur at IR and lower frequencies, where extinction is not sufficient and there is no lower energy photon field to bulk-Comptonize to IR frequencies. While the Klein-Nishina effect can flatten the slope at the high-energy end of the relativistic electron energy distribution (e.g., Bottcher et al. 1997), it can not produce the two distinct synchrotron peaks that we observe in 3C 454.3 (Fig. 14). We therefore focus on the scenario of two synchrotron emission regions.
Our observations of 3C 454.3 provide some clues to the relationship between the two (mid-IR and sub-mm) synchrotron emission peaks. First, that both peaks occur in both the high and low states, over a period of 6 months (and also in the 1996-1997 ISO observations) indicates that the jet structures that produce them must persist in both the flaring and quiescent jet states. They are not a transient phenomenon particular to the 2005 jet flaring episode. However, further monitoring at shorter timescales is crucial to elucidate the relationship between the two peaks. Second, the similar flux intensity (in νF ν ) and factor ∼ 40 increase in the flux of both synchrotron peaks indicates that they come from regions or components in the near-side jet that have similar Doppler beaming factors (δ ∼ 10 from VLBI radio observations; see §1).
The very dissimilar peak frequencies (2 × 10 11 vs. 8 × 10 13 Hz), along with the lack of correlated variability at short timescales between the two peaks ( Fig. 12) indicates that they are produced in distinct regions of the jet with different physical parameters. The primary mid-IR peak likely comes from a more compact region closer to the base of the jet than the sub-mm peak because of synchtrotron self-absorption. In the 2005 outburst, both regions appear to have responded to a major dis-turbance traveling along the jet. If the two regions are at different radii along the jet, then relativistic light travel time effects must shorten the apparent time between the disturbance at the first and second regions. The peak of the 2005 optical outburst occured sometime before 9 May 2005 (the start of a WEBT observing campaign). The 2005 July sub-mm flare occured sometime between 16 May and 3 June 2005 ∆t ∼ 7−25 days (> 0.5−2×10 6 s) after the peak of the optical outburst. Assuming the disturbance travelled down the jet at relativistic speed v, arriving first at the IR emission region and second at the sub-mm emission region, then we estimate the distance between the two regions to be ∆r = βc∆t/(1+z)
We consider two possible geometries for the two synchrotron emission regions (Fig. 15). The geometry suggested by Raiteri et al. (2009) for BL Lac is a helical jet, where the emission comes from two blobs moving at different angles to the line of sight, corresponding to different Doppler parameters. However, it is not clear in this picture why there would be two distinct synchrotron blobs rather than 3 or more blobs or a continuous distribution of emission along the helix. Rather, the existence of exactly two synchrotron peaks which are present in both the low and high states suggests a more permanent jet structure.
Another possible geometry for the two synchrotron emission peaks is that the IR peak comes from the base of the jet, where it becomes optically thin to mid-IR photons, while the sub-mm peak comes from a recollimation shock located further down the jet. Such a recollimation shock has been predicted with jet hydrodynamic simulations (Lind et al. 1989), where jet recollimation is caused by magnetic hoop stresses within the jet. This scenario has recently been invoked to explain the stationary (yet optically variable) HST 1 synchrotron emission knot observed in the jet of M87 (Nakamura, Garofolo, & Meier 2010, in preparation). Alternatively, the sub-mm synchrotron emission region may mark a location where the jet runs into an external density enhancement such as the broad-line region.
If two distinct synchrotron-emitting populations of particles are moving with different speeds and orientations to our line of sight, both populations would have to be folded into models that compute the contributions to the spectral energy distribution from synchrotron self-Compton emission, and external Compton emission from the accretion disk, broad line region, and corona, as has been done for a single broken-power-law-distribution population by Vercellone et al. (2009a). The resulting SEDs may show two peaks in the γ-ray band, corresponding to the two peaks in the infrared and sub-mm bands.

The Origin of mid-IR Flaring
Our observations may be further interpreted in the framework of the relativistic shock model (Marscher & Gear 1985). In this model, shocks are introduced at the base of a pre-existing relativistic jet by events in the accretion disk or supermassive black hole ergosphere. Shocks may be produced by an increase of the relativistic electron energy density, bulk flow speed, or changes in the magnetic field configuration. The electron density, energy density, and magnetic field will be enhanced at a shock front, leading to copious synchrotron and inverse Compton emission observed as a bright flare at radio-gamma frequencies.
According to the shock model, the peak of the mid-IR emission comes from a location along the relativistic jet where the shock becomes optically thin to synchrotron self-absorption. As the shock region progresses along the jet and expands, it will become optically thin at progressively lower frequency, and the synchrotron peak will move from the mid-IR into the far IR as it fades. The historical SED (Figure 14) shows that a synchrotron peak at ∼ 10 12 Hz (∼ 300 µm, rest) is more typical of 3C 454.3 in a non-flaring state.
As the shock propagates down the jet, it loses energy primarily through inverse Compton scattering. An increase in the ambient photon density field, e.g. flares from the accretion disk or close-encounters with broadline clouds would drain energy from the shock via Compton drag. This can not explain mid-IR flaring, though it could certainly cause flaring of Compton X-rays and γ-rays, as for 3C279 (Wehrle et al. 1998).
The shock may be re-energized as it runs into inhomogeneities in the jet or surrounding interstellar medium and bulk kinetic energy is converted to electron kinetic energy and magnetic fields. This could explain why we see large-amplitude mid-IR flaring with Spitzer two months after the peak of the 2005 May optical flare, after the original shock had propagated away from the energy source at the base of the jet.
Variability on longer time scales could also be caused by changes in the Doppler factor as the shock velocity changes its angle to the line of sight. As pointed out for BL Lac by Raiteri et al. (2009) and for 3C 454.3 by Villata et al. (2007), beaming in a curved jet varies along the jet and boosts the emission from various parts of the jet which are emitting at different frequencies, possibly causing changes in the observed flux and spectral index.

Shock Parameters for the IR-Optical Emission
Region We use simple arguments to derive rough estimates of the physical parameters of the IR-optical flaring synchrotron emission regions in 2005 July. The source doubled its MIR (6 µm) flux in 5 days from 2005 July 8-13 (2.7 days in the host galaxy rest frame), giving a variability brightness temperature of T var = 3.0 × 10 10 K, close to the canonical equipartition value of 5 × 10 10 K. The corresponding variability size is < 3.7 × 10 12 cm if the shock has a bulk Doppler parameter of δ ∼ 10. This is smaller than the gravitational radius of the supermassive black hole, R G = 1.5 × 10 13 M 8 cm.
Optical synchrotron emission may be produced when the magnetic field and Lorentz factor of the emitting elec-trons are in the range B ∼1-10 Gauss and γ ∼ 10 3 − 10 4 . We calculate the e-folding time for synchrotron cooling (t cooling = 2.4 × 10 9 γ −1 4 B −2 µG yr, where γ 4 = γ/10 4 and B is in units of microGauss) using Eq. 9.23 of Krolik (1999) for B=10 Gauss and γ = 10 4 and 10 3 of ∼13 minutes and ∼ 2 hours, respectively. Both of these values are consistent with variations on timescales of days and an upper limit of one day for the time delay between optical and infrared bands. The overall decline in optical flux on e-folding timescales of six months may be due to adiabatic expansion of a newly emitted or newly coalesced jet component. We find that: 1) There are no significant narrow or broad emission or absorption features in either the low or high-resolution mid-IR spectra of the high state or low resolution spectrum of the low state, consistent with a dominant nonthermal (synchrotron) emission mechanism during both states.
2) The mid-IR continuum emission is ∼ 30 times brighter than during previous ISO observations. The mid-IR flux is highly variable, decreasing by a factor of of 20-40 across the 5-35 µm band between July 2005 and December 2006, indicating Doppler boosted synchrotron emission from the jet.
3) There are two variable synchrotron emission peaks, one in the sub-mm, and one in the mid-IR. The lack of correlated variability between the peaks on 1-30 day time scales and a lag of 7-25 days between the sub-mm and optical outbursts indicates two separate synchrotron emission regions, separated by ∼ 0.07 − 5 pc.
4) The frequency of the primary synchrotron peak varies from < 10 13 Hz to 4×10 13 Hz, moving through the Spitzer IRS bandpass. The peak frequency is higher in outburst than during the low state, consistent with injection of high-energy electrons in a jet shock.
5) The 6-24 µm spectral slope becomes harder during flares and the peak of the SED moves to higher frequency. This may indicate that the shock is being re-energized as a new jet component runs into inhomogeneities in a slower, preexisting jet or surrounding ISM. 6) No time delay between optical and infrared light curves was observed; any time delay is smaller than 0.5 day. However, there are significant differences between the flaring behavior in the optical and infrared bands. Some flares have larger amplitude in the mid-IR, while others have larger amplitude in the optical.
7) The near-IR -optical spectral slope varies dramatically on the timescale of days, perhaps indicating injection of high energy electrons in jet shocks.
This work is based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407. Optical observations at Palomar Observatory were made possible through the NASA Space Interferometry Mission Preparatory Science Program for the Key Project on Active Galactic Nuclei. Support for this work was provided by NASA through awards issued by JPL/Caltech. The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica. This work is partly based on data taken and assembled by the WEBT collaboration and stored in the WEBT archive t the Osservatorio Astronomico di Torino-INAF (http://www.oato.inaf.it/blazars/webt/). We are grateful to Karl Gordon (Arizona) and Jason Surace (Spitzer Science Center) for assistance in processing MIPS and IRAC data, and to Booth Hartley for customizing IPAC Skyview for use on Intel Macs. We are grateful to Massimo Villata and Claudia Raiteri for providing WEBT data and for helpful discussions, and to Bill Keel (U. of Alabama) and Marco Chiaberge (STScI) for working with us on the Hubble images.                         1 and 15), with low resolution spectra (black lines) from the same epochs overlaid. Background has been subtracted for the low resolution spectra but not for the high resolution spectra (for which we had no separate background measurements), hence, there is an offset between the two spectra. Top panel: SH; bottom panel: LH.      (Villata et al. 2006) are included with the Colgate and Palomar data to improve the sampling during the 2005 June-July Spitzer observation perod. The mm-wave variations do not appear to be correlated with the mid-IR variations, suggesting that emission in the two bands comes from separate regions in the jet. c) Cross-correlation of IRAC 8.0 µm and R-band light curves (solid line). The autocorrelation curves for IRAC 8.0 µm (dash), and R-band (dot-dash) are shown for comparison. The IR and optical emission are strongly correlated at zero lag.  (Villata et al. 2007). For comparison, we show the rise in the SED (at ∼ 10 15 Hz) attributed by Raiteri et al. (2008a) to the Big Blue Bump; B,U, W1, and M2 data points are from the Optical Monitor on XMM, taken on 2006 Dec 18-19, (Table 3 of Raiteri et al. (2007) ).  Table 3. V, R, and I photometry for 2006 Dec 20 are from WEBT data provided by M. Villata (Raiteri et al. 2007;Villata et al. 2007). Both sets have been corrected for extinction. Chandra X-ray data for the 2007 minimum are are indicated by the error box. The scattered small black dots are radio and optical photometry from the NED database, gathered from 1979-1995. The mean Richards et al. (2006) QSO SED (dashed line) and the Rieke et al. 10 13 L ⊙ ULIRG template (dotted line) are scaled and overplotted for comparison. Neither of these SEDs is compatible with the shape of the IR bump in the low-state 3C 454.3 SED, indicating a dominant nonthermal contribution from the jet. Note the presence of two synchrotron peaks, at IR and sub-mm wavelengths, in both the low and high states. The IR-optical peak is produced closer to the base of the jet, while the submm peak is produced further out. B) Inhomogeneous jet model. Two radially separated synchrotron emission regions give rise to the IR and submm synchrotron emission peaks. The IR emission region comes from the base of the jet, while the sub-mm emission region comes from a shock (possibly a jet recollimation shock) at a radius of ∼ 0.07 − 5 pc. The direction to the observer's line of sight is marked "L. O. S."